ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 Photometry of the GCs in NGC 1700 has been discussed by 2.., Photometry of the GCs in NGC 1700 has been discussed by \scite{whitmore97}.3 Here we reanalyse the cata. ancl use the results to aid with GC selection in our Ixeck data.," Here we reanalyse the data, and use the results to aid with GC selection in our Keck data."4 We performed the photometry in 2 pixel. radius apertures with (?).. applving the aperture correction given by ο and the calibration and. transformation to Johnson magnitudes following ?..," We performed the photometry in 2 pixel radius apertures with \cite{bertin96}, applying the aperture correction given by \scite{whitmore97} and the calibration and transformation to Johnson magnitudes following \scite{holtzman95}. ."5 From the initial object list of 383 sources detected. by we applied the following selection criteria: The magnitude range was restricted to be 21.5«V 26.5. the [aint limit chosen to avoid. introducing colour bias effects," From the initial object list of 383 sources detected by we applied the following selection criteria: The magnitude range was restricted to be $21.5<V<26.5$ , the faint limit chosen to avoid introducing colour bias effects"6stars: quasars are brighter relatively than stars at both short wavelengths (the UVX method) and long wavelengths (the IxX method).,stars; quasars are brighter relatively than stars at both short wavelengths (the UVX method) and long wavelengths (the KX method).7 In this section. we review the motivation [for surveys that can detect reddened quasars., In this section we review the motivation for surveys that can detect reddened quasars.8 In 822 we describe the method., In 2 we describe the method.9 Then. in 833 we consider the feasibility. of undertaking a large KN survey for dampec να galaxies and gravitational lenses using the planned. UINIR fielcl nearinfrared. camera.," Then, in 3 we consider the feasibility of undertaking a large KX survey for damped $\alpha$ galaxies and gravitational lenses using the planned UKIRT wide--field near–infrared camera."10 Samples of distant quasars that are bright at. optical wavelengths have proven valuable in a number of areas in extraGalactic astronomy ancl cosmology., Samples of distant quasars that are bright at optical wavelengths have proven valuable in a number of areas in extra–Galactic astronomy and cosmology.11 Here. we focus on the use of quasars for ii) Aharting the history of star formation with lookback time through surveys. [or dampec Lya (DLA) galaxies. (Pei and. Fall. 1995). and ii) determining the geometry of the Universe. [rom the nmieasurement of the [frequency of gravitational lensing (Fukugita. Putamase and Ixasai. 1990. Turner 1990).," Here we focus on the use of quasars for i) charting the history of star formation with lookback time through surveys for damped $\alpha$ (DLA) galaxies (Pei and Fall 1995), and ii) determining the geometry of the Universe from the measurement of the frequency of gravitational lensing (Fukugita, Futamase and Kasai 1990, Turner 1990)."12 The elfects of dust in the intervening galaxies are crucial to the interpretation of such studies., The effects of dust in the intervening galaxies are crucial to the interpretation of such studies.13 Surveys for quasars at optical wavelengths. are very susceptible to. extinction. because the observed. passband corresponds to. the. rest-frame ultraviolet at moderate redshift and bevond.," Surveys for quasars at optical wavelengths are very susceptible to extinction, because the observed passband corresponds to the rest-frame ultraviolet at moderate redshift and beyond."14 Utilising nearinfrared wavelengths. produces a significant. improvement., Utilising near--infrared wavelengths produces a significant improvement.15 For example. a highredshift quasar. z3. behind a DLA system at 2=2.5 with (restframe) (0V)=0.2 mag would suller 2.8 mag extinction in the observedframe Dband but only 0.5 mag at Ix. (IEExeept where otherwise noted all extinction calculations assume an LMClike extinction curve below 0.335 rest.frame (Llowarth 1983). the Mathis (1990) curve at longer wavelengths. anc fy= 3.1).," For example, a high–redshift quasar, $z \sim 3$, behind a DLA system at $z=2.5$ with (rest–frame) $E(B-V)= 0.2\,$ mag would suffer $2.8\,$ mag extinction in the observed–frame B--band but only $0.5\,$ mag at K. (Except where otherwise noted all extinction calculations assume an LMC–like extinction curve below $0.33\mu$ rest–frame (Howarth 1983), the Mathis (1990) curve at longer wavelengths, and $R_V=3.1$ )."16 ‘To measure the history of gas consumption by star formation in the Universe. quasars are used as background sources of light ancl their spectra are searched for DLA absorption lines caused bv intervening clouds with a high column density of neutral hydrogen.," To measure the history of gas consumption by star formation in the Universe, quasars are used as background sources of light and their spectra are searched for DLA absorption lines caused by intervening clouds with a high column density of neutral hydrogen."17 Bright quasars are needed to provide sullicienthy high signaltonoise ratio spectra to measure the absorption line properties of the intervening gas., Bright quasars are needed to provide sufficiently high signal–to–noise ratio spectra to measure the absorption line properties of the intervening gas.18 The frequeney and. column. densities of the DLA absorption lines vield a cireet measurement of the cosmic density of neutral gas. Quy. (Wolfe 1987).," The frequency and column densities of the DLA absorption lines yield a direct measurement of the cosmic density of neutral gas, $\Omega_{\rm19HI}$, (Wolfe 1987)."20 Since clusty DLA absorbers can be missed because the quasar will be. dimmed. by extinction below the Ilux limit chosen for the spectroscopic observations. the procedure is only valid.," Since dusty DLA absorbers can be missed because the quasar will be dimmed by extinction below the flux limit chosen for the spectroscopic observations, the procedure is only valid."21"accound. The measured decline of Quy, with time. corrected for the effects of dust. can be related cürectiv to the history of star formation (under assumptions about the exchange of gas between the neutral and ionised gas phases)"," The measured decline of $\Omega_{\rm HI}$ with time, corrected for the effects of dust, can be related directly to the history of star formation (under assumptions about the exchange of gas between the neutral and ionised gas phases)."22 Pei and Fall (1995) have made the most detailed and complete analysis of this problem., Pei and Fall (1995) have made the most detailed and complete analysis of this problem.23 To account for the effects of dust. they made. the simplifving assumption that all absorbers at the same redshift have the same dustto.gas ratio., To account for the effects of dust they made the simplifying assumption that all absorbers at the same redshift have the same dust–to–gas ratio.24 By including in their treatment a termi specifying the mean metallicitv.. Z(z). of the DLA absorbers they successfully computed: models. of cosmic star. formation and chemical enrichment that account for the. recdshift dependence of the measured 2g. while allowing. in a selfconsistent manner. for the selection ellect of the increasing obscuration due to dust as star formation progresses.," By including in their treatment a term specifying the mean metallicity, $Z(z)$, of the DLA absorbers they successfully computed models of cosmic star formation and chemical enrichment that account for the redshift dependence of the measured $\Omega_{\rm HI}$, while allowing, in a self--consistent manner, for the selection effect of the increasing obscuration due to dust as star formation progresses."25 The outputs of the calculations are the evolution with redshift of the true Og. the rate of star formation and the metallicity of the gas.," The outputs of the calculations are the evolution with redshift of the true $\Omega_{\rm HI}$ , the rate of star formation and the metallicity of the gas."26 Pei and Fall confirmed that a mocest range in the dustgas ratio in DLAs at any redshift does not greatly alter their results., Pei and Fall confirmed that a modest range in the dust--to--gas ratio in DLAs at any redshift does not greatly alter their results.27 Phe actual range in the dusttogas ratio is cillicult to measure. because any clusty DLAs will have been missed.," The actual range in the dust–to–gas ratio is difficult to measure, because any dusty DLAs will have been missed."28 Boissé et al (1998) provide a discussion of this issue ancl present evidence for a bias against clusty DLAs in the current samples., Boissé et al (1998) provide a discussion of this issue and present evidence for a bias against dusty DLAs in the current samples.29 Therefore. it appears at least. possible that the actual range in dusttogas ratios is large. ancl that many cust. DLAs have evaded. the census. which would mean that our picture of the history of gas consumption is incorrect.," Therefore, it appears at least possible that the actual range in dust–to–gas ratios is large, and that many dusty DLAs have evaded the census, which would mean that our picture of the history of gas consumption is incorrect."30 A large sample of DLAs selected from spectra of a sample of quasars Εανlimited in the Ix.band would provide the data fora proper treatment of the ellects of dust., A large sample of DLAs selected from spectra of a sample of quasars flux–limited in the K–band would provide the data for a proper treatment of the effects of dust.31" Μαιστία ct al (1990) and Turner (1990) show that for Lat geometries as predicted by inflation. Qu,|Ox=1. the probability that a quasar is gravitationally lensed is more than an order of magnitude greater for largo values of the cosmological constant A (O4~ 1) than for small values (Q4~ 0)."," Fukugita et al (1990) and Turner (1990) show that for flat geometries as predicted by inflation, $\Omega_{matter}+\Omega_{\Lambda}=1$, the probability that a quasar is gravitationally lensed is more than an order of magnitude greater for large values of the cosmological constant $\Lambda$ $\Omega_{\Lambda}\sim 1$ ) than for small values $\Omega_{\Lambda}\sim 0$ )."32 Therefore. counting the fraction of gravitational lenses in a sample of quasars is in principle a powerful method for establishing the global geometry of the Universe.," Therefore, counting the fraction of gravitational lenses in a sample of quasars is in principle a powerful method for establishing the global geometry of the Universe."33 In practice the calculations necessary are quite involved (e.g. Alaoz and Rix 1993. Ixochanek 1996) but the conclusion of most analvses has been that the statistics o£ [ensec quasars and AGN are inconsistent with a large value for the cosmological constant.," In practice the calculations necessary are quite involved (e.g. Maoz and Rix 1993, Kochanek 1996) but the conclusion of most analyses has been that the statistics of lensed quasars and AGN are inconsistent with a large value for the cosmological constant."34 For example Paleo. Kochanck anc Alunoz (1998) provide the 2e limit Qy<0.62.," For example Falco, Kochanek and Munoz (1998) provide the $2\sigma$ limit $\Omega_{\Lambda}<0.62$."35 However. in a recent stuck. Chiba and. Yoshii (1999) found a bes fit value O4.~0.7.," However, in a recent study, Chiba and Yoshii (1999) found a best fit value $\Omega_{\Lambda}\sim 0.7$."36 The cillerence between these results is largely explained. by uncertainties in the properties of the lensing galaxies (velocity dispersions. space densities). which probably prohibit a definitive answer until they are better established.," The difference between these results is largely explained by uncertainties in the properties of the lensing galaxies (velocity dispersions, space densities), which probably prohibit a definitive answer until they are better established."37 Another potentially significant source of uncertainty is our lack of knowledge of the evolution. of massive galaxies. in particular their merging history.," Another potentially significant source of uncertainty is our lack of knowledge of the evolution of massive galaxies, in particular their merging history."38 All calculations agreeο that massive earlytvtypepe ggalaxies dominate the lensing Crosssection., All calculations agree that massive early–type galaxies dominate the lensing cross–section.39 Locally. most earlytype ealaxies contain litte dust so the existence of a number of highly reddened. lensed. quasars (e.g. MGO414|0534. Lawrence et al 1995) is something of a puzzle.," Locally, most early–type galaxies contain little dust so the existence of a number of highly reddened, lensed, quasars (e.g. MG0414+0534, Lawrence et al 1995) is something of a puzzle."40 Phe reddening may be intrinsic to he quasar but. reddening due todust, The reddening may be intrinsic to the quasar but reddening due todust41Figure 12. shows the profiles of the (init=1 discs. averaged over the final 13 ORPs.,"Figure \ref{fig:q1} shows the profiles of the $q_{\rm init}=1$ discs, averaged over the final 13 ORPs."42" The initial stellar masses are AZ,= 0.5A/.. M,=JAL.. and M,=2M.."," The initial stellar masses are $M_* = 0.5 M_\odot$ , $M_* = 1 M_\odot$, and $M_* = 2 M_\odot$."43 The discs grow hotter with increasing dise mass (with a flatter temperature profile). while maintaining a similar surface density profile.," The discs grow hotter with increasing disc mass (with a flatter temperature profile), while maintaining a similar surface density profile."44 This results in the higher dise mass simulations obtaining a flatter aspect ratio (top right panel)., This results in the higher disc mass simulations obtaining a flatter aspect ratio (top right panel).45 Figure 13.> shows that all the discs have similar qualitative a profiles. with τρια being different to what would be expected if the local approximation were appropriate (ecu).," Figure \ref{fig:q1alpha} shows that all the discs have similar qualitative $\alpha$ profiles, with $\alpha_{\rm total}$ being different to what would be expected if the local approximation were appropriate $\alpha_{\rm cool}$ )."46 The value of the enhancement appears to increase with increasing dise mass. showing that while q dictates whether or not a dise deviates from the local approximation. the dise mass Λι controls the strength of this deviation (through its influence on X and ultimately the dise thickness).," The value of the enhancement appears to increase with increasing disc mass, showing that while $q$ dictates whether or not a disc deviates from the local approximation, the disc mass $M_{\rm d}$ controls the strength of this deviation (through its influence on $\Sigma$ and ultimately the disc thickness)."47 All three discs have aspect ratios in excess of 0.1 for most of their radial extent. again consistent with previous predictions for non-locality (2)..," All three discs have aspect ratios in excess of 0.1 for most of their radial extent, again consistent with previous predictions for non-locality \citep{Lodato_and_Rice_04}."48 Figure |4 also shows that the quasi-steady approximation also appears to be violated (Figure 143)., Figure \ref{fig:TQ_q1} also shows that the quasi-steady approximation also appears to be violated (Figure \ref{fig:TQ_q1}) ).49 The temperature fluctuates at values of  and higher. with similar fluctuations in Q.," The temperature fluctuates at values of $\sim$ and higher, with similar fluctuations in $Q$."50 The non-local transport fraction (bottom right panel in Figure 13))in all three cases Is | or larger showing that the transport is very non-local., The non-local transport fraction (bottom right panel in Figure \ref{fig:q1alpha}) )in all three cases is $\sim 1$ or larger showing that the transport is very non-local.51the expanded SCE potential.,the expanded SCF potential.52 The initial conditions. and the models used. to. represent the primary and. satellite are summarised in Table L1., The initial conditions and the models used to represent the primary and satellite are summarised in Table \ref{modtbl}.53 Curve (a) in the Figure is a simple analytical approximation of the sinking rate based on the Chandrasekhar dynamical friction. formula (Binney&Tremaine 1987)., Curve (a) in the Figure is a simple analytical approximation of the sinking rate based on the Chandrasekhar dynamical friction formula \cite{bt}. .54. We approximate f(0) —constant. and assume instantaneously circular orbits for the satellite.," We approximate $f(0)=$ constant, and assume instantaneously circular orbits for the satellite."55 This calculation carries with it the assumption of a fixed primary and ids valid only in the region of large separations. +. low circular. velocity. no satellite mass loss. and is to be disregarded at small r.," This calculation carries with it the assumption of a fixed primary and is valid only in the region of large separations, $r$, low circular velocity, no satellite mass loss, and is to be disregarded at small $r$."56 As one might expect when the SCE expansion is truneated at low orders curves (d) and (0)] dynamical friction is poorly modelled., As one might expect when the SCF expansion is truncated at low orders [curves (d) and (e)] dynamical friction is poorly modelled.57 Increasing the truncation order curves (b) and (0) with a corresponding increase in the resolution of the SCE system causes the Tree system to respond realistically to perturbations in the SCL density field., Increasing the truncation order [curves (b) and (c)] with a corresponding increase in the resolution of the SCF system causes the Tree system to respond realistically to perturbations in the SCF density field.58 With the SCE expansion truncated at nm=16 and {=6 and taking a single point mass for the satellite one tree. particle) dynamical (rietion is modelled: well (see curve c]., With the SCF expansion truncated at $n=16$ and $l=6$ and taking a single point mass for the satellite one tree particle) dynamical friction is modelled well (see curve [c]).59 We find these sinking times are consistent with the results obtained by Hernquist and Weinberg (1989) for fully self-consistent simulations using a pure tree code and concur with their lindings and of White (1983)] that when the response of the primary is included in the numerical calculations the sinking time increases by a factor of around two., We find these sinking times are consistent with the results obtained by Hernquist and Weinberg (1989) for fully self-consistent simulations using a pure tree code and concur with their findings [and of White (1983)] that when the response of the primary is included in the numerical calculations the sinking time increases by a factor of around two.60 Decay of the orbit of a single point mass means in simple terms that orbital energy is transferred to dynamical heating of the particles in the primary., Decay of the orbit of a single point mass means in simple terms that orbital energy is transferred to dynamical heating of the particles in the primary.61 For an extended: satellite one composed of INusus bodies). orbital energy can also be transferred to the heating of the satellite. thereby increasing the decay rate and shortening the decay. period.," For an extended satellite one composed of $N_{satellite}$ bodies), orbital energy can also be transferred to the heating of the satellite, thereby increasing the decay rate and shortening the decay period."62 This interesting result is demonstrated in curve (b). and certainly warrants greater detailed exploration.," This interesting result is demonstrated in curve (b), and certainly warrants greater detailed exploration."63 The result has been clearly alluded to in Weinberg (1989) in terms of coupling in the analytic linear theory which he describes., The result has been clearly alluded to in Weinberg (1989) in terms of coupling in the analytic linear theory which he describes.64 'The extended satellite has many weak resonances which can couple at smaller radii. which the point mass does not.," The extended satellite has many weak resonances which can couple at smaller radii, which the point mass does not."65 Finally we note that the sinking satellite loses much of its mass curing its descent., 	 Finally we note that the sinking satellite loses much of its mass during its descent.66 For the example traced by curve (b) in Figure(13)). of the Tree particles become unbound from tree svstem at r~7. ancl are unbound alr~2.," For the example traced by curve (b) in \ref{sink}) ), of the Tree particles become unbound from tree system at $r\sim 7$, and are unbound at $r\sim 2$."67 The peculiar behaviour of the curves oncehaving reached ro Lis due to the practical disruption of the satellite., The peculiar behaviour of the curves oncehaving reached $r\sim 1$ is due to the practical disruption of the satellite.68the Milkv. Way disk. and (hat clump giants are reliable distance indicators and (racers of the old and intermediate age populations.,"the Milky Way disk, and that clump giants are reliable distance indicators and tracers of the old and intermediate age populations."69 However. the Besancoon Galactic model al.2003) wilh a scale-leneth of 2.4 kpc. including a flaring and warped outer disk. with a disk eutoff at 2=14 kpc seems to reproduce well the observed distance distributions (see Fig.4)).," However, the Besançoon Galactic model \citep{2003A&A...409..523R} with a scale-length of 2.4 kpc, including a flaring and warped outer disk, with a disk cutoff at $R=14$ kpc seems to reproduce well the observed distance distributions (see \ref{model}) )."70 Data and model show the same features. will a sharp increase. (he peak around 5 kpc and (he exponential decrease.," Data and model show the same features, with a sharp increase, the peak around 5 kpc and the exponential decrease."71 The edege is detected in the data and model at the expected place within the errors |d=10.1 kpe (corresponding to R=13.9 kpe). and d=9.5 kpe (R=13.4 kpc). respectively].," The edge is detected in the data and model at the expected place within the errors $d=10.1$ kpc (corresponding to $R=13.9$ kpc), and $d=9.5$ kpc $R=13.4$ kpc), respectively]."72" A [ew studies have discussed the edge of the Milky Way towards the anticentre region previously,", A few studies have discussed the edge of the Milky Way towards the anticentre region previously.73 The old Galactic disk apparently does not extend bevond 14 kpc on the basis of optical V. vs. (D—V) colour-magnitude diagrams and star counts (Robinetal.1992).. which does not contradict the presence of other more distant. voung stars.," The old Galactic disk apparently does not extend beyond 14 kpc on the basis of optical $V$ vs. $(B-V)$ colour-magnitude diagrams and star counts \citep{1992ApJ...400L..25R}, which does not contradict the presence of other more distant young stars."74 Models of the COBE-DIRBE infrared emission maps vield (hat the edge of the disk is at 4 kpe from the Sun (assuming{ο=8.5kpe.Freudenreich.1996)..," Models of the COBE-DIRBE infrared emission maps yield that the edge of the disk is at 4 kpc from the Sun \citep[assuming75 $R_o=8.5$~kpc,][]{1996ApJ...468..663F}."76 DENIS star counts reveal the cut off of the stellar disk al a Galactocentric distance of 1542 kpe (Ruphyetal.1996).., DENIS star counts reveal the cut off of the stellar disk at a Galactocentric distance of $15\pm2$ kpc \citep{1996A&A...313L..21R}.77 In contrast. star counts from 2\LASS reveal no radial disk trimeation at 14 kpe 2006)..," In contrast, star counts from 2MASS reveal no radial disk truncation at 14 kpc \citep{2006A&A...451..515M}."78 Again. the depth of 2\IASS is not enough to reach large distances. which is possible with the VVV and UIXIDSS-GPS survey that reach 3-4 magnitudes fainter.," Again, the depth of 2MASS is not enough to reach large distances, which is possible with the VVV and UKIDSS-GPS survey that reach 3-4 magnitudes fainter."79 More recently. it was found that the 2ALASS star counts are best fit if the external disk is truncated at 12-14 kpe (Revléetal.2009).. while early A-type stars in the anticentre from the IPILAS survey reveal an exponential disk out to 13 kpe. with a steeper decline bevond that distance (Saleetal.2010)..," More recently, it was found that the 2MASS star counts are best fit if the external disk is truncated at 12-14 kpc \citep{2009A&A...495..819R}, while early A-type stars in the anticentre from the IPHAS survey reveal an exponential disk out to 13 kpc, with a steeper decline beyond that distance \citep{2010MNRAS.402..713S}."80 Most of the previous evidence for cutoffs in the stellar distribution has been acquired at the anticentre fields., Most of the previous evidence for cutoffs in the stellar distribution has been acquired at the anticentre fields.81 The present deep exploration of different fields in three ealactic quacdrants finds consistent results., The present deep exploration of different fields in three galactic quadrants finds consistent results.82 Even though we interpret this termination of the chimp giant distribution as a türuncation. we nole that a break in the slope of the distribution would also be consistent with (he cata.," Even though we interpret this termination of the clump giant distribution as a truncation, we note that a break in the slope of the distribution would also be consistent with the data."83 It is difficult to measure (he sharpness of the cutolL. because the number of stars drop rapidly with distance. bul as the surveys progress. we will have more fields in order to explore this Issue.," It is difficult to measure the sharpness of the cutoff, because the number of stars drop rapidly with distance, but as the surveys progress, we will have more fields in order to explore this issue."84 This does not mean that one could not find other voung stellar sources bevond the distance that we measure., This does not mean that one could not find other young stellar sources beyond the distance that we measure.85 There are stars detected bevond the edge of the old stellar disk., There are stars detected beyond the edge of the old stellar disk.86 For example. (here is a population of distant voung stars at 2=20 kpc. in two fields located," For example, there is a population of distant young stars at $R=20$ kpc, in two fields located"87in the Y and J-bands.,in the Y and J-bands.88 Chioromethans. acetvlene and hydrocarbons used here show small contributions.," Chloromethans, acetylene and hydrocarbons used here show small contributions."89 Also the absorption gas cell proposed by D’Amatoetal.(2008) covers the II and Ix-bands with ~ 2200 lines but again there is a lack for Y and J-bands calibration with few absorptions., Also the absorption gas cell proposed by \citet{2008SPIE.7014E.126D} covers the H and K-bands with $\sim$ 200 lines but again there is a lack for Y and J-bands calibration with few absorptions.90 Obtaining deeper absorptions would imply (o increase the length of the gas cell. but woulel also imply a deeper absorption for the IE and We have presented new results of gas mixtures [or wavelength calibration echelle spectrographis in (he near infrared.," Obtaining deeper absorptions would imply to increase the length of the gas cell, but would also imply a deeper absorption for the H and We have presented new results of gas mixtures for wavelength calibration echelle spectrographs in the near infrared."91 We have worked on different. gas cells including several new gases., We have worked on different gas cells including several new gases.92 The working method and the properties of the gas cells have been described., The working method and the properties of the gas cells have been described.93 We have obtained several mixtures and we have presented a compact and manageable gas cell which covers (he widest wavelength range to date in the II and Ix-bands. with a potentially hieh number of lines than for currently available gas cells. stable in time scales of months under abmospheric temperature conditions of our laboratory. aud which can be useful for hieh precision radial velocity. measurements.," We have obtained several mixtures and we have presented a compact and manageable gas cell which covers the widest wavelength range to date in the H and K-bands, with a potentially high number of lines than for currently available gas cells, stable in time scales of months under atmospheric temperature conditions of our laboratory, and which can be useful for high precision radial velocity measurements."94 We work on the improvement of the gas cell using different partial pressures of the individual gases in order to solve the pressure broadening of some absorption bands., We work on the improvement of the gas cell using different partial pressures of the individual gases in order to solve the pressure broadening of some absorption bands.95 Some of these gases have been recently tested with real observations and obtained promising results with few m/s accuracy., Some of these gases have been recently tested with real observations and obtained promising results with few m/s accuracy.96" Such gas cells can be of interest for several new eeneration high resolution near infrared spectrographs under development (NAIIUAL. GIANO. ΡΗΝΟ, SPIROU. This work has been supported bv the Spanish Ministerio de Eduacionn ν΄ Ciencia through erant. AYA2004-08271-CO1."," Such gas cells can be of interest for several new generation high resolution near infrared spectrographs under development (NAHUAL, GIANO, PRVS, SPIROU, This work has been supported by the Spanish Ministerio de Eduaciónn y Ciencia through grant AYA2004-08271-C01."97 Work on writing this paper was developed while LV was a visitor at the Centro de Estudios de Físsica del Cosmos de Aragónn. whose hospitality is acknowledged gratefully.," Work on writing this paper was developed while LV was a visitor at the Centro de Estudios de Físsica del Cosmos de Aragónn, whose hospitality is acknowledged gratefully."98"For the time-averaged meridional (low we use the same profile as ?.. Le. The two remaining parameters of the fIux transport model are (he horizontal and racial diffusiviües. gj, and 5.","For the time-averaged meridional flow we use the same profile as \cite{van_Ballegooijen98}, i.e. The two remaining parameters of the flux transport model are the horizontal and radial diffusivities, $\eta_{H}$ and $\eta_r$."99 The results of various attempls to measure 7j; [rom observation are summarized in Table 6.2 of ?.., The results of various attempts to measure $\eta_{H}$ from observation are summarized in Table 6.2 of \cite{Schrijver_book}.100 The values obtained from cross-correlation and methods [all in the range 100—300 kin?s !.., The values obtained from cross-correlation and object-tracking methods fall in the range $100-300$ $^2$ $^{-1}$.101 We have used jjj=250 Κιν [ον our reference value in Section 4: this value lies within the range of the observations. but we also consider the effect of varving it in Section 5.," We have used $\eta_{H}=250$ $^{2}$ $^{-1}$ for our reference value in Section 4; this value lies within the range of the observations, but we also consider the effect of varying it in Section 5."102 Aluch less is known about 5., Much less is known about $\eta_r$.103 This term was introduced bv ο to account for the 3D racial diffusion of the magnetic field and to obtain regularly reversing polar fields for eveles of varving amplitude in (he absence of variations of the meridional flow., This term was introduced by \cite{Baumann06} to account for the 3D radial diffusion of the magnetic field and to obtain regularly reversing polar fields for cycles of varying amplitude in the absence of variations of the meridional flow.104 Its physical motivation is Chat the σας magnetic field is three-dimensional aud thus has more modes of decay than are captured by (he two-dimensional surface diffusion., Its physical motivation is that the Sun's magnetic field is three-dimensional and thus has more modes of decay than are captured by the two-dimensional surface diffusion.105" We find here that the results with pj,=0 match the observations well (including having the polar fields reverse each cycle) when we include the observed till angle variations.", We find here that the results with $\eta_r=0$ match the observations well (including having the polar fields reverse each cycle) when we include the observed tilt angle variations.106 We thus take jj.=0 as our reference value and consider other values in Section 5., We thus take $\eta_r=0$ as our reference value and consider other values in Section 5.107 For the source term 5(À.6./) in Equation 1 we follow ? and ? and consider new flux {ο emerge in the form of of opposite polarity patches.," For the source term $S(\lambda,\phi,t)$ in Equation \ref{eqn:SFT} we follow \cite{van_Ballegooijen98} and \cite{Baumann04} and consider new flux to emerge in the form of of opposite polarity patches."108 The positive-polaritv patch is centered on latitude A and longitude © . the negative patch at (A.).," The positive-polarity patch is centered on latitude $\lambda_+$ and longitude $\phi_+$ , the negative patch at $(\lambda_-,\phi_-)$."109 The field of each new bipole is givenbv 5-2B8—D. with where ο.(λ.ὁ) are (he heliocentric angles between (λ.ὁ) and (AL.62). respectively and Ao=9(A.ó) is the separation between the two polarities. and 9=4 is the size of the individual polaritv patches.," The field of each new bipole is given by $B=B^+-B^-$ with where $\beta_{\pm}(\lambda,\phi)$ are the heliocentric angles between $(\lambda,\phi)$ and $(\lambda_{\pm},\phi_{\pm})$, respectively and $\Delta \beta=\beta_+(\lambda_-,\phi_-)$ is the separation between the two polarities, and $\delta=4^\circ$ is the size of the individual polarity patches."110 For the purposes of comparing the flix transport simulations with observations it is necessary to connect S closely to the actual observations., For the purposes of comparing the flux transport simulations with observations it is necessary to connect $S$ closely to the actual observations.111 We use sunspot group areas andl locations corresponding to their (me of maxinuun area from (based on the Greenwich photoheliographic maps from 1874 to 1976 and USAF/NOAA SOON data thereafter) as proxies for emerging fIux., We use sunspot group areas and locations corresponding to their time of maximum area from (based on the Greenwich photoheliographic maps from 1874 to 1976 and USAF/NOAA SOON data thereafter) as proxies for emerging flux.112 The Greenwich/USAF/NOAA record contains the locations and areas of sunspots groups. but no magneticpolarity information.," The Greenwich/USAF/NOAA record contains the locations and areas of sunspots groups, but no magneticpolarity information."113 We use (he location and areas (ο construct bipolar, We use the location and areas to construct bipolar114The black hole candidate 33394. was discovered by Alarkertetal.(1973) with the satellite anc was soon noted for its similarity. in X-ray behaviour to the classical black hole candidate XN1 (Alarkertetal. 1973: Macjimaetal. 1984: Dolanetal. 1987)).,The black hole candidate 339–4 was discovered by \scite{markert73} with the satellite and was soon noted for its similarity in X-ray behaviour to the classical black hole candidate X–1 \pcite{markert73}; \pcite{maejima84}; \pcite{dolan87}) ).115 The source exhibits aperiodie and quasi-periodic modulations on ime scales spanning milliseconds to vears and over a wide range of wavelengths., The source exhibits aperiodic and quasi-periodic modulations on time scales spanning milliseconds to years and over a wide range of wavelengths.116 Unlike tvpical soft. X-ray transients (SNT8) which are barely detected in their quiescent state. 33394 spends most of the time in the so-called: X-ray ονπαρ state (LS) like a faint persistent source.," Unlike typical soft X-ray transients (SXTs) which are barely detected in their quiescent state, 339–4 spends most of the time in the so-called X-ray low/hard state (LS) like a faint persistent source."117 During he LS. the energy spectrum can be described as a power-aw with photon index of ~ 1.52 (see Appendix for full ist. of references).," During the LS, the energy spectrum can be described as a power-law with photon index of $\sim$ 1.5–2 (see Appendix for full list of references)."118 Lt changes to the high/soft state (118) occasionally (c.g. Maejiniaetal.1984:Bellonict 1999)) which are similar to the outbursts of transient. sources.," It changes to the high/soft state (HS) occasionally (e.g. \pcite{maejima84,belloni99}) ) which are similar to the outbursts of transient sources."119 During the LIS. it becomes. brighter (in. the 210. keV icd) by a factor of ~ 5LOO (see e.g. Ixongetal.2000.. ancl references therein) ancl exhibits an ultra-soft. spectral component plus a steeper power-law component.," During the HS, it becomes brighter (in the 2–10 keV band) by a factor of $\sim$ 5–100 (see e.g. \pcite{kong00}, and references therein) and exhibits an ultra-soft spectral component plus a steeper power-law component."120" The Low"" and ""high states referred to here are also called the ""hard? and ‘soft’ states so as to rellect the spectral behaviour.", The `low' and `high' states referred to here are also called the `hard' and `soft' states so as to reflect the spectral behaviour.121 llowever. the luminosity of the ‘low’ state can sometimes be higher than the ‘high’ state (see Appencix) and hence it can be confusing to define the ‘state’ solely bv the intensity of the N-ravs.," However, the luminosity of the `low' state can sometimes be higher than the `high' state (see Appendix) and hence it can be confusing to define the `state' solely by the intensity of the X-rays."122 Therefore. a knowledge of the X-ray spectra are essential to distinguish cillerent states accurately (sec ‘Table 1 of Ixongctal.9000 for the spectral ancl temporal properties of dillerent states)," Therefore, a knowledge of the X-ray spectra are essential to distinguish different states accurately (see Table 1 of \pcite{kong00} for the spectral and temporal properties of different states)."123 1n addition to the LS and. HIS. 33394. exhibits a very high state (VIIS: Mivamotoctal. 1991)). with a ügher X-ray luminosity than in the LIS (bv a factor. of ~ 3).," In addition to the LS and HS, 339–4 exhibits a very high state (VHS; \pcite{miyamoto91}) ) with a higher X-ray luminosity than in the HS (by a factor of $\sim$ 3)."124 Note that the VIIS. also occurs in several other xack hole soft. X-ray transients (BIISNEs): 11124.683 (Ebiswwaetal.1994).. 11550564 (Sobezak1999) and XTEJJ1748288 (Revnivtisey.TrudolvubovDorozdin 2000).," Note that the VHS also occurs in several other black hole soft X-ray transients (BHSXTs): 1124–683 \cite{ebisawa94}, 1550–564 \cite{sobczak99} and J1748–288 \cite{revnivtsev00}."125.. An intermediate state (LS) in e.g. (Aléndez&vanderIxlis1997). and 11630.47 (Dietersal.2000). has also been reported. ancl its spectral and iming properties are similar to the VIIS but with a much ower luminositv., An intermediate state (IS) in e.g. 339--4 \cite{mendez97} and 1630–47 \cite{dieters00} has also been reported and its spectral and timing properties are similar to the VHS but with a much lower luminosity.126 Finally. every now and then 33394 enters an “oll” state (see Markertetal. Ό," Finally, every now and then 339–4 enters an `off' state (see \pcite{markert73}; ;"127 Finally. every now and then 33394 enters an “oll” state (see Markertetal. Ότ," Finally, every now and then 339–4 enters an `off' state (see \pcite{markert73}; ;"128 Finally. every now and then 33394 enters an “oll” state (see Markertetal. Ότο," Finally, every now and then 339–4 enters an `off' state (see \pcite{markert73}; ;"129the disk experiences significant changes.,the disk experiences significant changes.130 Asiueutioned in 2.1.. the inuer portion of the disk may become unstable? 1 when the accretion rate is relatively high.," As mentioned in \ref{ssec:spin}, the inner portion of the disk may become unstable \cite{le74,kro98}, when the accretion rate is relatively high."131 Once this occurs. the disk mielt become fragmented.," Once this occurs, the disk might become fragmented."132 The eravitomaguetic precession of the fragiucuts is probably only weakly damped. so it might be capable of producing the observed QPOs.," The gravitomagnetic precession of the fragments is probably only weakly damped, so it might be capable of producing the observed QPOs."133 A recent theoretical developiueit involves the discovery of a class of high-frequency eravitomagnetic mode sthat are only weakly damped. in addition to the strougly-cdamped low-frequeow oues P.," A recent theoretical development involves the discovery of a class of high-frequency gravitomagnetic modes that are only weakly damped, in addition to the strongly-damped low-frequency ones \cite{ml98}."134 These high-frequeney modes are very localized. spiral corrugations o: the inner portion of the disk., These high-frequency modes are very localized spiral corrugations of the inner portion of the disk.135 They are perhaps most relevant to the οSCLved QPOs. because X-ray ΙΕ frou DIIBCs is ikev to oriellate ver vecose to the ceutral black hole.," They are perhaps most relevant to the observed QPOs, because X-ray emission from BHBCs is likely to originate very close to the central black hole."136 However. it remains to © see1 whoether the LOes can be excited easily aud whether they are capahle of luodulatiie N-raw enaission 10," However, it remains to be seen whether the modes can be excited easily and whether they are capable of modulating X-ray emission \cite{ml98}."137 It shoul ITE pointed out that wo critical assuniptfions are mace in the study: sna ilt arele and (pseudo) Newtonian poteitial., It should be pointed out that two critical assumptions are made in the study: small tilt angle and (pseudo) Newtonian potential.138 The former simplifies the οςuatiol bv ienorme the non-luear aspect of the problem. which might affect the «wuping11 time scales of eravitomaenoetic 1nodes m a significant wav: non-li1041 Calculaious are required to shed light ou this ise.," The former simplifies the equation by ignoring the non-linear aspect of the problem, which might affect the damping time scales of gravitomagnetic modes in a significant way; non-linear calculations are required to shed light on this issue."139 The latter migit still represen a good approximation for svstenis that contain a relatively slowly rotating neutron star. but is certainly invalid for rapidly rotating black holes in which the stable QPOs are observed.," The latter might still represent a good approximation for systems that contain a relatively slowly rotating neutron star, but is certainly invalid for rapidly rotating black holes in which the stable QPOs are observed."140 TLOSC asstuptious. therefore. tend [9] limit the scope of the applicability of the results.," These assumptions, therefore, tend to limit the scope of the applicability of the results."141 At present. if would seeu premature to draw any definitive conclusions about confirming or rejecting the eravitomaguetic origin of the stable QPOs in nicroquasars.," At present, it would seem premature to draw any definitive conclusions about confirming or rejecting the gravitomagnetic origin of the stable QPOs in microquasars."142 The eravitoniagjetfie precession of accreted matter oilv provides a natural frequency for thi ΣΩΡΟ»., The gravitomagnetic precession of accreted matter only provides a natural frequency for the QPOs.143 To actually see them. certain physical processes are required to procuce X-ray modulations.," To actually see them, certain physical processes are required to produce X-ray modulations."144 The processes are οο] wulsaown at present., The processes are entirely unknown at present.145" Possiülities include (by no means exclusivevk (1) variation in the ταν enütti warea of the disk. due to precession: (2) Doppler or gravitational shift of phot neiere. due to the Ixepleriau motion of sclf-illi1ninating chunips or ""hot spots” iu the accretion disk: (3) oscillation that uodulates the emissiou from the disk: Hard CL) occultation of a highly compact X-ray cinitting region by disk fraegmenus or the inner part of the disk itself."," Possibilities include (by no means exclusively): (1) variation in the X-ray emitting area of the disk, due to precession; (2) Doppler or gravitational shift of photon energy, due to the Keplerian motion of self-illuminating clumps or “hot spots” in the accretion disk; (3) oscillation that modulates the emission from the disk; and (4) occultation of a highly compact X-ray emitting region by disk fragments or the inner part of the disk itself."146 The transicut nature of the QPOs secus to imply tiat the first possibility. if viable at a1. iust be only a small effect.," The transient nature of the QPOs seems to imply that the first possibility, if viable at all, must be only a small effect."147 Απ processes that involve discrete, Any processes that involve discrete148for having. while in outburst. a magnetosphere and. of course. an accurately measured spin period.,"for having, while in outburst, a magnetosphere and, of course, an accurately measured spin period."149 These characteristics make the source very well suited for testing the predictions of models for the quiescent emission in which the presence of a sizeable magnetic field plays a crucial role (Stella et al., These characteristics make the source very well suited for testing the predictions of models for the quiescent emission in which the presence of a sizeable magnetic field plays a crucial role (Stella et al.150 1994)., 1994).151 SAX J1808.4-3658 was detected in quiescence (Stella et al., SAX J1808.4–3658 was detected in quiescence (Stella et al.152 2000: Dotani et al., 2000; Dotani et al.153 2001; Wijnands et al., 2001; Wijnands et al.154 2001) though with large uncertainties., 2001) though with large uncertainties.155 Here we report on an XMM-Newton observation of SAX J1808.4-3658 in quiescence. the first to detect the source with a good signal to noise ratio.," Here we report on an XMM-Newton observation of SAX J1808.4–3658 in quiescence, the first to detect the source with a good signal to noise ratio."156 SAX J1808.41—3658 was observed on March 24 2001 with XMM-Newton EPIC. consisting of two metal oxide semiconductor (MOS) cameras (Watson et al.," SAX J1808.4–3658 was observed on March 24 2001 with XMM-Newton EPIC, consisting of two metal oxide semiconductor (MOS) cameras (Watson et al."157 2000) and one pn camera (Strüdder et al., 2000) and one pn camera (Strüdder et al.158 2000)., 2000).159 Medium filters were used for the MOS cameras and the thin filter for the pn camera., Medium filters were used for the MOS cameras and the thin filter for the pn camera.160 The pn camera was operated in timing mode in order to search for pulsations in case the source were sufficiently bright., The pn camera was operated in timing mode in order to search for pulsations in case the source were sufficiently bright.161 SAX J1808.4—3658 turned out to be faint (see below). preventing meaningful searches for pulsed emission.," SAX J1808.4–3658 turned out to be faint (see below), preventing meaningful searches for pulsed emission."162 We do not discuss the pn camera data in the following and concentrate on data from MOS cameras that were operated in full frame mode with a read-out time of 2.6 s. We extracted the event file starting from the raw data using the Standard Analysis System (SAS) version 5.3.0., We do not discuss the pn camera data in the following and concentrate on data from MOS cameras that were operated in full frame mode with a read-out time of 2.6 s. We extracted the event file starting from the raw data using the Standard Analysis System (SAS) version 5.3.0.163 Data were manually screened to remove any remaining bright pixels or hot columns., Data were manually screened to remove any remaining bright pixels or hot columns.164 Periods m which the background was high because of soft proton flares were excluded using an intensity filter: we rejected all events accumulated in the external CCDs 7) when the total count rate exceeded 15 counts in 100 s in the 10-12.4 keV band for the MOSI and MOS2. independently.," Periods in which the background was high because of soft proton flares were excluded using an intensity filter: we rejected all events accumulated in the external CCDs (2--7) when the total count rate exceeded 15 counts in 100 s in the 10–12.4 keV band for the MOS1 and MOS2, independently."165 We obtained a net exposure time of 33.2 ks and 34.4 ks for MOS] and MOS2. respectively (this is slightly different from what reported in Campana et al.," We obtained a net exposure time of 33.2 ks and 34.4 ks for MOS1 and MOS2, respectively (this is slightly different from what reported in Campana et al."166 2002. due to the use of standard processed files with SAS 5.0.3).," 2002, due to the use of standard processed files with SAS 5.0.3)."167 Spectra were accumulated separately for MOSI and MOS2., Spectra were accumulated separately for MOS1 and MOS2.168 Event grades higher than 12 were filtered out., Event grades higher than 12 were filtered out.169 Photons were grouped within to the nominal MOS resolution of 15 eV. Previous observations aiming at determining the position of SAX J1808.4-3658 were performed with BeppoSAX (Stella et al., Photons were grouped within to the nominal MOS resolution of 15 eV. Previous observations aiming at determining the position of SAX J1808.4–3658 were performed with BeppoSAX (Stella et al.170 2000; Wijnands et al., 2000; Wijnands et al.171 2002b) and ASCA (Dotani. Asal Wijnands 2000: see Table 1). but due to the limited. angular resolution of these satellites the resulting picture was quite confused.," 2002b) and ASCA (Dotani, Asai Wijnands 2000; see Table 1), but due to the limited angular resolution of these satellites the resulting picture was quite confused."172 In particular it was unclear whether the source revealed at à low flux of ~101°ergs!cn7 and with large positional uncertainties was indeed SAX J1808.4—3658 (see Fig., In particular it was unclear whether the source revealed at a low flux of $\sim 10^{-13}\ergs\cmdue$ and with large positional uncertainties was indeed SAX J1808.4–3658 (see Fig.173 | and Tab., 1 and Tab.174 1)., 1).175 In fact. Wijnands et al. (," In fact, Wijnands et al. ("1762002b) claimed that SAX J1808.4-3658 was only detected in the ASCA data of September 1999. while neither BeppoSAX observations detected the millisecond pulsar but rather revealed a nearby source. SAX J1808.6-3658.,"2002b) claimed that SAX J1808.4–3658 was only detected in the ASCA data of September 1999, while neither BeppoSAX observations detected the millisecond pulsar but rather revealed a nearby source, SAX J1808.6–3658."177 The deep image we obtained with EPIC allowed us to clarify this confused picture by revealing several faint sources in the region., The deep image we obtained with EPIC allowed us to clarify this confused picture by revealing several faint sources in the region.178 Source detection was obtained following the prescriptions of Baldi et al. (, Source detection was obtained following the prescriptions of Baldi et al. (1792002).,2002).180 This revealed 20 sources in the central part (12« 9) of the field of view., This revealed 20 sources in the central part $12'\times9'$ ) of the field of view.181 The brightest one was clearly coincident with the radio position of SAX J1808.6-3658., The brightest one was clearly coincident with the radio position of SAX J1808.6–3658.182 A few other sources were contained in the error regions of sources previously seen with ASCA and in the longer BeppoSAX observation (see Fig., A few other sources were contained in the error regions of sources previously seen with ASCA and in the longer BeppoSAX observation (see Fig.183 1)., 1).184" A faint source is visible about 2’ East of SAX J1808.4-3658 (RA(J2000); 18"" 089 367.6: DEC(2000):236* 58' 01"").", A faint source is visible about $2'$ East of SAX J1808.4–3658 (RA(J2000): $^{\rm h}$ $^{\rm m}$ $^{\rm s}$ .6; $\deg$ $58'$ $01''$ ).185 This is within the error region of the source detected by Wijnands et al. (, This is within the error region of the source detected by Wijnands et al. (1862002b) in the March 2000 BeppoSAX observation.,2002b) in the March 2000 BeppoSAX observation.187 If it had the same luminosity that we measure with EPIC (a count rate a factor ~5 lower than that of SAX J1808.4-3658). it was probably too faint to be detected by BeppoSAX. but a small contamination from this souree could perhaps explain why all the error regions for SAX J1808.4—3658 derived with low angular resolution instruments are systematically centered to the East of the radio position.," If it had the same luminosity that we measure with EPIC (a count rate a factor $\sim 5$ lower than that of SAX J1808.4–3658), it was probably too faint to be detected by BeppoSAX, but a small contamination from this source could perhaps explain why all the error regions for SAX J1808.4–3658 derived with low angular resolution instruments are systematically centered to the East of the radio position."188" A further bias is provided by a 2.9/ to the brighter source East (RAQO2000): 18! 083 427.2: DEC(J2000):-36° Sa’ 11"") which has a comparable flux to SAX J1808.4-3658 and went undetected in previous observations (Wijnands et al.", A further bias is provided by a $2.9'$ to the brighter source East (RA(J2000): $^{\rm h}$ $^{\rm m}$ $^{\rm s}$ .2; $\deg$ $58'$ $44''$ ) which has a comparable flux to SAX J1808.4–3658 and went undetected in previous observations (Wijnands et al.189 2002b)., 2002b).190 Furthermore. given the source density revealed in this field by XMM-Newton. the probability of finding by chance at least one source within the BeppoSAX 90% error circle of SAX J1808.6-3658 is TOU.," Furthermore, given the source density revealed in this field by XMM-Newton, the probability of finding by chance at least one source within the BeppoSAX $90\%$ error circle of SAX J1808.6–3658 is $\sim 70\%$."191 It appears to be unlikely that this faint source was much brighter and at the same time SAX J1808.4—3658 much fainter Just during the BeppoSAX observations and conclude that SAX J1808.4—3658 was the source detected both in March 1999 and March 2000., It appears to be unlikely that this faint source was much brighter and at the same time SAX J1808.4–3658 much fainter just during the BeppoSAX observations and conclude that SAX J1808.4–3658 was the source detected both in March 1999 and March 2000.192" In total we collected 91 and 77 counts with the MOSI and MOS2 detectors. respectively. within a 10"". radius circle centered on the position of SAX J1808.4—-3658."," In total we collected 91 and 77 counts with the MOS1 and MOS2 detectors, respectively, within a $10''$ radius circle centered on the position of SAX J1808.4–3658."193 This corresponds to about 55.—GO% of the source flux for energies between 1-5 keV (extending the extraction region would just decrease the signal to noise ratio. due to the background and source faintness).," This corresponds to about $55-60\%$ of the source flux for energies between 1–5 keV (extending the extraction region would just decrease the signal to noise ratio, due to the background and source faintness)."194" The background was extracted from an annular region with inner and outer radii of 0.5"" and 2.5 (excluding contaminating sources). respectively."," The background was extracted from an annular region with inner and outer radii of $0.5'$ and $2.5'$ (excluding contaminating sources), respectively."195 Within. the SAX J1808.1—3658 extraction region. about 15.—204 of the counts are to be attributed to the local background.," Within the SAX J1808.4–3658 extraction region, about $15-20\%$ of the counts are to be attributed to the local background."196 The MOSI and MOS2 spectrà were rebinnec to contain at least IS photons per channel., The MOS1 and MOS2 spectra were rebinned to contain at least 15 photons per channel.197 The spectral analysis owas carried ο with. XSPEC v11.1.0., The spectral analysis was carried out with XSPEC v.11.1.0.198 We adopted the latest on-axis standard response matrices. anc ο) and correct the fluxes for the fraction of photons collected within the extractior radius., We adopted the latest on-axis standard response matrices and ) and correct the fluxes for the fraction of photons collected within the extraction radius.199 The spectral analysis was carried out in the 0.3—7 keV energy range (Kirsch private communication and 2002)., The spectral analysis was carried out in the 0.3–7 keV energy range (Kirsch private communication and 2002).200 Giver that the cross calibration between MOS] and MOS? agrees within 5'4 (Kirsch 2002) we did not include a constant for the two MOS instruments., Given that the cross calibration between MOS1 and MOS2 agrees within $5\%$ (Kirsch 2002) we did not include a constant for the two MOS instruments.201 The statistics was relatively poor and the spectrum could be well fit by an absorbed power law (or a bremsstrahlung) model (see Table 2 and Fig., The statistics was relatively poor and the spectrum could be well fit by an absorbed power law (or a bremsstrahlung) model (see Table 2 and Fig.202 2)., 2).203 A power law model gave a better fit Qa= 1.1) with D—15/53 (0% confidence level) and a column density of Ny=0.3107em ος1&1074e.? 90% cL.," A power law model gave a better fit $\chi^2_{\rm red}=1.1$ ) with $\Gamma=1.5^{+0.2}_{-0.3}$ $90\%$ confidence level) and a column density of $N_H=0.3\times10^{21}\cmdue$ $<1.1\times10^{21}\cmdue$ $90\%$ c.l.,"204 we used here the Wisconsin cross sections in the model)., we used here the Wisconsin cross sections in the model).205 We note that a model with the column density fixed at the value derived from the source data in outburst (e.g. Gilfanov et al., We note that a model with the column density fixed at the value derived from the source data in outburst (e.g. Gilfanov et al.206" 1998) and consistent with the galactic value of 1.3«10?!«in.7. provides also a good fit with P=Lsc0.3 and 42,,=1.2 (with a null hypothesis probability of 20%)."," 1998) and consistent with the galactic value of $1.3\times 10^{21}\cmdue$, provides also a good fit with $\Gamma=1.8\pm0.3$ and $\chi^2_{\rm red}=1.2$ (with a null hypothesis probability of $20\%$ )."207 The unabsorbed 0.5-10 keV luminosity of SAX 11808.4—3658 was 5«10°!ores1 for the power law model and a (revised) distance of 2.5 kpe (in't, The unabsorbed 0.5–10 keV luminosity of SAX J1808.4–3658 was $5\times10^{31}\ergs$ for the power law model and a (revised) distance of 2.5 kpc (in't208Little is known of the parental molecular clouds of the IRDCs.,Little is known of the parental molecular clouds of the IRDCs.209 kinematic distance estimates., inematic distance estimates.210"where 0X(r) is the amplitude varying slowly with radius. and |11.5In(r,/0.21)+my] represents (he phase varving rapidly with radius. Equation (9)) and the condition dtr)>0 on the initial phase imply that the maximum values of the perturbed density in equation (10)) coincide with the positions of all the planets (Fig.","where $\delta \Sigma (r)$ is the amplitude varying slowly with radius, and $\left[ 11.5 \ln (r_{\mathrm{n}}/0.21)211+ m \varphi \right]$ represents the phase varying rapidly with radius, Equation \ref{eq:titbod2}) ) and the condition $\delta \Sigma (r)>0$ on the initial phase imply that the maximum values of the perturbed density in equation \ref{eq:cos}) ) coincide with the positions of all the planets (Fig."212 2aa), \ref{fig:positions}a a).213 Interestingly. and this is the central part of our theory. the TB rule leq. (10))]satisfies," Interestingly, and this is the central part of our theory, the TB rule [eq. \ref{eq:cos}) )]"214 ihe eonditions of the WIND wave withthe effective TB radial wavenumber dlnhkuy/dinr= O(1). and Agr>1 f," the conditions of the WKB wave withthe effective TB radial wavenumber $\mathrm{d} \ln k_{\mathrm{eff}}/\mathrm{d}\ln r =O(1)$ , and $k_215{\mathrm{eff}}r \gg 1$ [cf."216el. ο., eq. \ref{eq:number})217"two most commonly used de-dispersion scripts, the one in PRESTO and the one in SIGPROC.","two most commonly used de-dispersion scripts, the one in PRESTO and the one in SIGPROC."218 A fake file was generated containing a 600-second observation centered at 153 MHz with a bandwidth of 6.24 MHz split into 1024 channels and having a sampling rate of 165 ms (containing a total of about 3.6x10° samples)., A fake file was generated containing a 600-second observation centered at 153 MHz with a bandwidth of 6.24 MHz split into 1024 channels and having a sampling rate of 165 ms (containing a total of about $3.6\times10^6$ samples).219" This was run through the three software suites for avalue.. For the GPU code and PRESTO only the actual de-dispersion part was timed, whilst SIGPROC also loads from file in the innermost loop so the timing contains some file I/O time as well."," This was run through the three software suites for a. For the GPU code and PRESTO only the actual de-dispersion part was timed, whilst SIGPROC also loads from file in the innermost loop so the timing contains some file I/O time as well."220 The timings are listed in table 3.., The timings are listed in table \ref{bruteComparisonTable}.221" The algorithm used to perform subband de-dispersion (both steps) is almost identical to the one used in brute-force de-dispersion, so the scaling tests were not repeated."," The algorithm used to perform subband de-dispersion (both steps) is almost identical to the one used in brute-force de-dispersion, so the scaling tests were not repeated."222 Figure 8 depicts a comparison plot between the two algorithms for various de-dispersion parameters.," Figure \ref{subbandComparisonFigure}223 depicts a comparison plot between the two algorithms for various de-dispersion parameters."224 The speed-up factor depends on optimal parameter combinations as well as the number of subbands and number of nominal DM values in the DM range employed for subband de-dispersion., The speed-up factor depends on optimal parameter combinations as well as the number of subbands and number of nominal DM values in the DM range employed for subband de-dispersion.225" The speed-up factor decreases linearly with increasing number of nominal DM values for a particular range, since more work needs to be done in the first algorithm step (although the same number of DM values are processed, the first step will generally be more intensive since the data has not been reduced yet)."," The speed-up factor decreases linearly with increasing number of nominal DM values for a particular range, since more work needs to be done in the first algorithm step (although the same number of DM values are processed, the first step will generally be more intensive since the data has not been reduced yet)."226" The number of nominal DM values and the number of subbands depend on the amount of dispersion smearing permissible, where a large subband-DM step and few subbands result in a higher amount of smearing."," The number of nominal DM values and the number of subbands depend on the amount of dispersion smearing permissible, where a large subband-DM step and few subbands result in a higher amount of smearing."227 These values should be fine-tuned to acquire the best balance between S/N and processing speed., These values should be fine-tuned to acquire the best balance between S/N and processing speed.228" The GPU subband de-dispersion implementation was also compared with PRESTO's prepsubband script, on which the algorithm is based."," The GPU subband de-dispersion implementation was also compared with PRESTO's prepsubband script, on which the algorithm is based."229 A fake file for a single-beam 60-second observation at 300 MHz with a bandwidth of 16 MHz and 1024 channels was created., A fake file for a single-beam 60-second observation at 300 MHz with a bandwidth of 16 MHz and 1024 channels was created.230 The plan used for the test is listed in table 4.., The plan used for the test is listed in table \ref{subbandPlanTestTable}.231 The time taken for the GPU code and PRESTO to process the entire file isrespectively., The time taken for the GPU code and PRESTO to process the entire file is.232". Again, this indicates that the GPU code is about two orders of magnitude faster than the CPU implementation, in this case 84x faster."," Again, this indicates that the GPU code is about two orders of magnitude faster than the CPU implementation, in this case $\times$ faster."233" For this test, PRESTO was run in single-thread mode."," For this test, PRESTO was run in single-thread mode."234 The performance boost obtained from GPUs make them an ideal candidate for use in real-time systems., The performance boost obtained from GPUs make them an ideal candidate for use in real-time systems.235 An off-the-shelf server with a high-end CUDA-enabled graphics card has enough power to de-disperse thousands of DM values in real-time (depending on telescope parameters)., An off-the-shelf server with a high-end CUDA-enabled graphics card has enough power to de-disperse thousands of DM values in real-time (depending on telescope parameters).236" Additional features are required for such a system, such as a way to read and interpret incoming telescope data, further channelisation and buffering between the input stream and de-dispersion buffers."," Additional features are required for such a system, such as a way to read and interpret incoming telescope data, further channelisation and buffering between the input stream and de-dispersion buffers."237" As a proof of concept, the GPU code was extended to include a channeliser (a simple FFT using the NVIDIA"," As a proof of concept, the GPU code was extended to include a channeliser (a simple FFT using the NVIDIA"238introduce a “blind” PCM method with Lstep foreerouud cleaning.,introduce a “blind” PCM method with 4-step foreground cleaning.239 We firstly eroup into pairs from the maps which have stroug cross correlations in phase., We firstly group into pairs from the maps which have strong cross correlations in phase.240" We them iu step b use a simple TOIT wmininization of the variance of the colmbination between differcut pairs of IKδν πας,", We then in step ${\bf b}$ use a simple TOH minimization of the variance of the combination between different pairs of K–W maps.241 In the step ο we propose a pliase-cleauiug filter for cach derived nap from the coustraimt of munimal variance of plase difference., In the step ${\bf c}$ we propose a phase-cleaning filter for each derived map from the constraint of minimal variance of phase difference.242 Finally. we propose a MIEN-MAX filter in order o obtain the PCM. map.," Finally, we propose a MIN-MAX filter in order to obtain the PCM map."243 Our CMD signal (the PCM) las the power spectrum in an excellent agreement with he best fit cosinological model. but the pliases are slightly differcut from those of the ILC imap.," Our CMB signal (the PCM) has the power spectrum in an excellent agreement with the best fit cosmological model, but the phases are slightly different from those of the ILC map."244 We also coufini that. after phase cleaning. the PCM as essentially no correlation iu phases with their cross-correlation with the foregrounds.," We also confirm that, after phase cleaning, the PCM has essentially no correlation in phases with their cross-correlation with the foregrounds."245 We would like to poiut out that the ΡΟΝΤ method can be used with different conrponeut separation methods such as tle Maxiunuu Entropy Method (MEM) or another methods. which cau be used an initial step before the phase cleaning.," We would like to point out that the PCM method can be used with different component separation methods such as the Maximum Entropy Method (MEM) or another methods, which can be used an initial step before the phase cleaning."246 Such eoncralization on the PCM method is in progress., Such generalization on the PCM method is in progress.247 Using the PCM method we have found. corresponding corrections of the foreground maps for cach EKW baud. which is useful for the data analysis for the wpcomine auission.," Using the PCM method we have found corresponding corrections of the foreground maps for each K–W band, which is useful for the data analysis for the upcoming mission."248 The power spectra of the como oreeround maps from our PCM subtraction are slightly differcut from the foreeround maps., The power spectra of the common foreground maps from our PCM subtraction are slightly different from the foreground maps.249 Moreover. the PCM does not show any serious contamination from the »oiut-like source residues mainly thauks to the MEN-MAX filter.," Moreover, the PCM does not show any serious contamination from the point-like source residues mainly thanks to the MIN-MAX filter."250 The PCM. CAB signal. as it follows frou comparison )etween our CAD andWAZA maps has a power smaller hen ILC imap.," The PCM CMB signal, as it follows from comparison between our CMB and maps has a power smaller then ILC map."251 We also confini that the PCAL displavs a low level of power at the quadrupole component. which also nanifest itself in the WALAPILC map aud the TOI nap (Teemark.deOliveiva-Costa&Taiilton2003:Oliveira-Costa.Tegmark.Zaldarriaga&Taimilton 2003).," We also confirm that the PCM displays a low level of power at the quadrupole component, which also manifest itself in the ILC map and the TOH map \citep{toh,oliveira}. ."252". llowever. the low level in €, is typical at the inultipole range 2«(x10."," However, the low level in $\Cl$ is typical at the multipole range $2 \le \ell \le 10$."253 In order to compare the power spectrum of the PCM with that of the best fit model and the TOT Wiener-fltered map. in Fig.l? we plot these power spectra with the error bars caused by the cosmic variance.," In order to compare the power spectrum of the PCM with that of the best fit model and the TOH Wiener-filtered map, in \ref{errorbars} we plot these power spectra with the error bars caused by the cosmic variance."254 As one can see from Fig.l? all deviations between the PCM and the TOT Wiencr-filtered map are within the CL., As one can see from \ref{errorbars} all deviations between the PCM and the TOH Wiener-filtered map are within the CL.255 Some of the multipole range. however. seenis to be quite peculiar.," Some of the multipole range, however, seems to be quite peculiar."256 Namely at 20<621 the power of both the PCM and the ΤΟΠ Wiener-filtered map πας some suppression.," Namely, at $20 \le \ell \le 24$ the power of both the PCM and the TOH Wiener-filtered map has some suppression."257 Does it indicate that we have a peculiar realization of a Claussian random CMD feld?, Does it indicate that we have a peculiar realization of a Gaussian random CMB field?258 Or we still ave the contauunatious of the foregrounds which iniuic as extra CMD power?, Or we still have the contaminations of the foregrounds which mimic as extra CMB power?259 All these issues can only be discussed when more data are available. particularly the future volarization observational data.," All these issues can only be discussed when more data are available, particularly the future polarization observational data."260 Iu this paper we concentrate ouly ou the low inultipole ranee of the CAIB (<50. at which the beam shape aud he instrumental noise are not crucial.," In this paper we concentrate only on the low multipole range of the CMB $\ell \le 50$, at which the beam shape and the instrumental noise are not crucial."261 The developinent of he PCM. imiethod and the results at the multipole rauge (>50 ds in preparation., The development of the PCM method and the results at the multipole range $\ell > 50$ is in preparation.262 This paper was supportec by Dammarks Gaundforskuinestoud through its support for the establishment of the Theoretical Astroplivsics Center., This paper was supported by Danmarks Grundforskningsfond through its support for the establishment of the Theoretical Astrophysics Center.263 We thank AL. Teemark et al., We thank M. Tegmark et al.264 for providing their processed naps., for providing their processed maps.265 We thank C. Efstathiou. A. Lasecuby. A. Doroshkevich. M. Demiauski aud P. R. Cliristeuseu for useful discussions.," We thank G. Efstathiou, A. Lasenby, A. Doroshkevich, M. Demianski and P. R. Christensen for useful discussions."266" We acknowledge the use of package (Córski.Hivoun&Wandelt1999) to produce ao, tom the data audthe use of the package (Doroshkevichetal.2003) for data analyses aud the whole-sky figures.", We acknowledge the use of package \citep{healpix} to produce $\alm$ from the data andthe use of the package \citep{glesp} for data analyses and the whole-sky figures.267giving andmasyr-!.,giving and.268. The Vizier database (Ochsenbeinetal.2000) lists five ground-based measurements that are in excellent agreement with the values., The Vizier database \citep{vizier} lists five ground-based measurements that are in excellent agreement with the values.269" Consequently, the current space velocity relative to the LSR iss-!."," Consequently, the current space velocity relative to the LSR is."270". This value transforms into a total Galactic rest frame velocity ofkms, slightly exceeding the Milky Way’s local escape velocity derived from the Galactic gravitational potential of(1991)."," This value transforms into a total Galactic rest frame velocity of, slightly exceeding the Milky Way's local escape velocity derived from the Galactic gravitational potential of."271 Hence a detailed kinematic study seemed worthwhile., Hence a detailed kinematic study seemed worthwhile.272" To do so, the Galactic potential of Allen&Santillan as well as the numerical code of Odenkirchen&Brosche(1992) was applied."," To do so, the Galactic potential of \citeauthor{alrmxaa22} as well as the numerical code of \citet{odan313} was applied."273 This allowed the stellar orbit to be traced back to the Galactic plane to investigate whether is à runaway star at all and to determine the flight time , This allowed the stellar orbit to be traced back to the Galactic plane to investigate whether is a runaway star at all and to determine the flight time $T_{\rm flight}$.274"Uncertainties were constrained via a Monte Carlo Thignt.method that simultaneously and independently varied the initial parameters (ie., the components of position and velocity), assuming for each a Gaussian distribution."," Uncertainties were constrained via a Monte Carlo method that simultaneously and independently varied the initial parameters (i.e., the components of position and velocity), assuming for each a Gaussian distribution."275 The following results are average or — when explicitly noted — median values of a sample of 0000. trajectories., The following results are average or – when explicitly noted – median values of a sample of 000 trajectories.276 Uncertainties are expressed by the standard deviation., Uncertainties are expressed by the standard deviation.277 The obtained travel time is shorter than the evolutionary timescale and, The obtained travel time is shorter than the evolutionary timescale and278 (Ixrisciunas et al., (Krisciunas et al.279 1995): these authors make a case that thefrequencies fj and. f» were constant from. 1992 January to 1994 February., 1995); these authors make a case that the $f_{1}$ and $f_{2}$ were constant from 1992 January to 1994 February.280 The large data set obtained by Zerbi et al., The large data set obtained by Zerbi et al.281 in 1994/5 allows us to wonder if the frequencies in the power spectrum. of the photometry of Xur are also variable., in 1994/5 allows us to wonder if the frequencies in the power spectrum of the photometry of Aur are also variable.282 In any case. it was found [rom the spectroscopic data of the 1993/4 season that the frequency lais the most pronounced one in. all three moment variations.," In any case, it was found from the spectroscopic data of the 1993/4 season that the frequency $^{-1}$ is the most pronounced one in all three moment variations."283 Lt was emphasized by WKrisciunas ct ((1995) that the dillerence in photometric and spectroscopic behavior of the frequencies f; ancl f» can be understood if fo would be due to a low-degree mode and fj to a higher-degree one., It was emphasized by Krisciunas et (1995) that the difference in photometric and spectroscopic behavior of the frequencies $f_1$ and $f_2$ can be understood if $f_2$ would be due to a low-degree mode and $f_1$ to a higher-degree one.284 This explanation is no longer appropriate if our identification is correct. since both modes have the same geometrical representation.," This explanation is no longer appropriate if our identification is correct, since both modes have the same geometrical representation."285 From the parameters given in (3)). we note that the amplitude of the racial part of the pulsation for fi isa factor 4 larger than the one for f». while its angular dependence is the same.," From the parameters given in \ref{opl}) ), we note that the amplitude of the radial part of the pulsation for $f_1$ is a factor 4 larger than the one for $f_2$, while its angular dependence is the same."286 Since the photometric variability is determined most of all by temperature variations and since the latter are determined by the radial part of the pulsation. it is quite understandable that the photometric variability is dominated by the mode with frequency. fi.," Since the photometric variability is determined most of all by temperature variations and since the latter are determined by the radial part of the pulsation, it is quite understandable that the photometric variability is dominated by the mode with frequency $f_1$."287 For the spectroscopic behavior. the radial displacement can completely be neglected in the case of large transverse amplitudes as we have here.," For the spectroscopic behavior, the radial displacement can completely be neglected in the case of large transverse amplitudes as we have here."288 The relative importance of roth modes for the line profiles could be the following., The relative importance of both modes for the line profiles could be the following.289 The displacement field £ correct up to first order in O7fy in the case of the two £—3 modes can easily be calculated. from the paper by erts Waelkens (1993).," The displacement field $\vec{\xi}$ correct up to first order in $\Omega/f_0$ in the case of the two $\ell=3$ modes can easily be calculated, from the paper by Aerts Waelkens (1993)."290 The toroidal corrections for the prograde mode (m=——1) and for the retrograde mode (m=|1) are dillerent., The toroidal corrections for the prograde mode $(m=-1)$ and for the retrograde mode $(m=+1)$ are different.291 Therefore. it is xossible that the total transverse velocity is larger for the mode with frequency fo than for the mode with frequcney fi. despite its smaller. spheroical amplitude ALL.," Therefore, it is possible that the total transverse velocity is larger for the mode with frequency $f_2$ than for the mode with frequency $f_1$, despite its smaller spheroidal amplitude $KA$."292 Quite uncierstandably. the line profiles are then mostly determined w fo.," Quite understandably, the line profiles are then mostly determined by $f_2$."293 In order to test this point. of view. we have explicitly calculated the velocities in the direction of the observer for roth modes. taking into account the toroidal corrections.," In order to test this point of view, we have explicitly calculated the velocities in the direction of the observer for both modes, taking into account the toroidal corrections."294 We indeed find that. although the retrograde mode has a our times smaller spheroidal zero-rotation amplitude. its otal velocity component in the line of sight becomes larger due to the toroidal correction terms.," We indeed find that, although the retrograde mode has a four times smaller spheroidal zero-rotation amplitude, its total velocity component in the line of sight becomes larger due to the toroidal correction terms."295 The maximum velocity reached by the retrograde mode is a factor 1.8 larger than he one of the prograde mode., The maximum velocity reached by the retrograde mode is a factor 1.8 larger than the one of the prograde mode.296 Lt thus seenis possible that the toroidal corrections tha appear due to influence of the rotation on a spheroicla mode introduce a different photometric and. spectroscopic behavior when both prograde anc retrograde. modes. are involved., It thus seems possible that the toroidal corrections that appear due to influence of the rotation on a spheroidal mode introduce a different photometric and spectroscopic behavior when both prograde and retrograde modes are involved.297 “To our knowledge. this is the first. detectec consequence of the presence of such toroidal corrections.," To our knowledge, this is the first detected consequence of the presence of such toroidal corrections."298 Our point of view could be further verified hy means of an analysis of a large set of high-resolution. high S/N spectra of a star pulsating in both prograde ancl retrograde mocdes ane having a large rotation frequency. hhaving fο)fy2120%.," Our point of view could be further verified by means of an analysis of a large set of high-resolution, high S/N spectra of a star pulsating in both prograde and retrograde modes and having a large rotation frequency, having $\Omega/f_0>20\%$."299 We are grateful to Dr. M. Mayor and the Observatoire de Geneve for having made the I-m Geneva telescope aud Coravel radial velocity spectrometer available to Dr. 1t. EF. Grillin for the data analyzed here., We are grateful to Dr. M. Mayor and the Observatoire de $\rm Gen\grave{e}ve$ for having made the 1-m Geneva telescope and Coravel radial velocity spectrometer available to Dr. R. F. Griffin for the data analyzed here.300 We thank Dr. Grillin for allowing himself to be convinced that the racial velocity of 3398 Aur might indeed be variable and worth looking into by means of systematic observations., We thank Dr. Griffin for allowing himself to be convinced that the radial velocity of 398 Aur might indeed be variable and worth looking into by means of systematic observations.301Lt is now widely accepted that most. normal galaxies have central. supermassive black holes (SMDBIIS) with Alby=10°PAL. Ixormendy.&Richiston1995:Mivoshiet.al.1995:Cillessenctal.2¢ )00)).,"It is now widely accepted that most normal galaxies have central supermassive black holes (SMBHs) with $M_{\rm BH} =10^{6\rm -9}M_{\odot}$ \citealt{KR95,Mi95,Gi09}) )."302 High resolution observations of nearby galaxies have revealed a close correlation between the mass of SAIBLIS and the bulge mass. or the velocity dispersion of the bulge (ος... Ixormendy&Richiston1995:ltix 2004)).," High resolution observations of nearby galaxies have revealed a close correlation between the mass of SMBHs and the bulge mass, or the velocity dispersion of the bulge (e.g., \citealt{KR95,FM00,MH03,HR04}) )."303 This implies that the formation of SAIBIIs is strongly coupled with t1e formation of galaxies., This implies that the formation of SMBHs is strongly coupled with the formation of galaxies.304 However. the formation and evoution of SMDBlIS is still an open question ancl a hot topic in astrophysics.," However, the formation and evolution of SMBHs is still an open question and a hot topic in astrophysics."305 The discovery of luminous quasars at redshift 26 shows that SAIBLIs of tve order of billions of solar masses exist at the end of the re-ionization epoch (e.g.. Fanοἱal.2001.2006:GotoWillottet2007. 2010)).," The discovery of luminous quasars at redshift $z >6$ shows that SMBHs of the order of billions of solar masses exist at the end of the re-ionization epoch (e.g., \citealt{Fa01,Fa06,Go06,Wi07,Wi10}) )."306 Considering Ededinegton-limitecl accretion. 26 quasars may not form reaclily.," Considering Eddington-limited accretion, $z > 6$ quasars may not form readily."307 Unless the seed. black hole mass is arge. eas acerction with low radiative clliciency and/or DII mergers significantly contribute to the growth of SALBLIs (sce Shapiro 20053).," Unless the seed black hole mass is large, gas accretion with low radiative efficiency and/or BH mergers significantly contribute to the growth of SMBHs (see \citealt{Sh05}) )."308 IE super-Idcdington accretion is possible. he growth timescale of 111” can be much shorter than the Eddington timescale (e.g. Ixawaguchietal.2004:Ohsugaal. 20053).," If super-Eddington accretion is possible, the growth timescale of BHs can be much shorter than the Eddington timescale (e.g., \citealt{Ka04,Oh05}) )."309 Thus. super-Eddington accretion might explain he formation of SMDBIIs.," Thus, super-Eddington accretion might explain the formation of SMBHs."310 Reeently. Kawakatu&Wacla(2009) predicted that super-IEdcdington accretion is required or the formation of +26 quasars with Aloz107AL.. oeause the final DII mass is ereathy suppressed. by star ormation in a cireumnuclear disc and AGN outllow based on he coevolution model of SMDLL growth and a cireumnuclear disc (Ixawalkatu&Wacla2008: see also Volonteri&Rees 20051).," Recently, \citet{KW09} predicted that super-Eddington accretion is required for the formation of $z>6$ quasars with $M_{\rm BH}\approx 31110^{9}M_{\odot}$, because the final BH mass is greatly suppressed by star formation in a circumnuclear disc and AGN outflow based on the coevolution model of SMBH growth and a circumnuclear disc \citealt{KW08}; see also \citealt{VR05}) )."312 The occurrence of super-Eddington accretion How has not vet been accurately verified. although this issue has been investigated since the 1970s (e.g.," The occurrence of super-Eddington accretion flow has not yet been accurately verified, although this issue has been investigated since the 1970s (e.g.,"313density profiles for the entire. metal-poor. aud metal-rich GC! populations. along with the PN profile.,"density profiles for the entire, metal-poor, and metal-rich GC populations, along with the PN profile."314 The populations were binned. following MaizApelliniz&Ubeda(2005).. into circular annuli of equal numbers of objects. providing the same statistical weight to each bin (although spatial biases may still affect the GC population in the outer regions along the major axis: see Fig. 3)).," The populations were binned, following \cite{maiz05}, into circular annuli of equal numbers of objects, providing the same statistical weight to each bin (although spatial biases may still affect the GC population in the outer regions along the major axis; see Fig. \ref{fig:gc_thetar}) )."315 In the inner 5 kpe of all tracer populations. incompleteness due to the obscuration of the dust lane is evident by the flattening of the surface density prolile.," In the inner 5 kpc of all tracer populations, incompleteness due to the obscuration of the dust lane is evident by the flattening of the surface density profile."316 The innermost objects were. therefore. excluded from the surface density profile fittings.," The innermost objects were, therefore, excluded from the surface density profile fittings."317 Outside οἱ 5 kpc. the surface densities lit well to power laws. leading to 5=3.65d:0.17. 3.49d0.34. 3.37+0.30. and 3.47£0.12 for the entire GC population. the metal-poor and metal-rich subpopulations of GCs. and the PNe in NGC! 5128. respectively.," Outside of 5 kpc, the surface densities fit well to power laws, leading to $\gamma = 3.65\pm0.17$, $3.49\pm0.34$, $3.37\pm0.30$, and $3.47\pm0.12$ for the entire GC population, the metal-poor and metal-rich subpopulations of GCs, and the PNe in NGC 5128, respectively."318 These are all very similar within their uncertainties., These are all very similar within their uncertainties.319 The similar kinematics we find between the metal-poor and metal-rich subpopulations of GCs in this study. strongly. justifies the combining of the two populations for the mass determination performed in Woodley(2006)., The similar kinematics we find between the metal-poor and metal-rich subpopulations of GCs in this study strongly justifies the combining of the two populations for the mass determination performed in \cite{w06}.320.. The GC population provides afolad mass estimate of (1.340.5)xLOY M. from 340 clusters ont to a projected radius of 50 kpe., The GC population provides a mass estimate of $(1.3\pm0.5) \times 10^{12}$ $M_{\odot}$ from 340 clusters out to a projected radius of 50 kpc.321 Removing the GCs in our sample with ο<300 km !|. which will remove all possible contamination from loreground stars. discussed in 3.1.. leads to a total mass of AL...," Removing the GCs in our sample with $v_r322\leq 300$ km $^{-1}$, which will remove all possible contamination from foreground stars, discussed in \ref{sec:velfield}, leads to a total mass of $(1.0\pm0.4) \times 10^{12}$ $M_{\odot}$."323 This mass agrees nicely with our mass determined [rom our entire GC sample., This mass agrees nicely with our mass determined from our entire GC sample.324 The PN population provides a total mass of (1.0220.2)x1072. M. [rom 780 PNe out to 90 kpe in projected radius. agreeing with the GC value within the uncertainty.," The PN population provides a total mass of $(1.0\pm0.2) \times 10^{12}$ $M_{\odot}$ from 780 PNe out to 90 kpc in projected radius, agreeing with the GC value within the uncertainty."325 We are also able to generate a mass profile of NGC 5123 from the total GC population and the PNe. shown in Figure 13..," We are also able to generate a mass profile of NGC 5128 from the total GC population and the PNe, shown in Figure \ref{fig:mass}."326 The tracer mass estimator determines (hefofal enclosed mass for NGC 5128 within the outermost radius of a given (tracer sample., The tracer mass estimator determines the enclosed mass for NGC 5128 within the outermost radius of a given tracer sample.327 It caleulates this total mass using a sample of objects defined within the radial range defined by the sample's inner and outermost radii., It calculates this total mass using a sample of objects defined within the radial range defined by the sample's inner and outermost radii.328 It is (therefore possible to. use a unique set of tracer objects. denoted by the radial bin range. listed in the [ist column of Tables 2--5.. to determine a mass prolile [rom independent mass estimates.," It is therefore possible to use a unique set of tracer objects, denoted by the radial bin range, listed in the first column of Tables \ref{tab:all_GC}- \ref{tab:PN}, to determine a mass profile from independent mass estimates."329 The independent binning. leads (o sample sizes in (he mass determination. in some cases generating higher uncertainties in the total enclosed mass.," The independent binning, leads to sample sizes in the mass determination, in some cases generating higher uncertainties in the total enclosed mass."330 The most certain mass is (he one determined from the Full sample of tracers., The most certain mass is the one determined from the full sample of tracers.331 In the mass determinations above. we have implicity. assumed isolropy [or (he velocity distributions.," In the mass determinations above, we have implicity assumed isotropy for the velocity distributions."332 But the possibility exists that the PNe (for example) mieht have radial anisotropy which would produce their gradually Falling σι(010) eurve., But the possibility exists that the PNe (for example) might have radial anisotropy which would produce their gradually falling $\sigma_v(R)$ curve.333 Replacing Equation 5 in the tracer, Replacing Equation \ref{eqn:C} in the tracer334"tha the ""canals are produced by beam depolarization.",that the `canals' are produced by beam depolarization.335 This perpeudiculuitv applies to all ‘canals’. respective of frequency band and is very convincing. especially because evervwlhere else the polarization vectors vary quite smoothly (if sienificautly!).," This perpendicularity applies to all `canals', irrespective of frequency band and is very convincing, especially because everywhere else the polarization vectors vary quite smoothly (if significantly!)."336 Beam depolarization creates ‘canals’ that are oue beam. wide. which is exactly what we observe.," Beam depolarization creates `canals' that are one beam wide, which is exactly what we observe."337 This nuplies that the 90° “jump? must occur on angular scales smaller than the beaunidth., This implies that the $^{\circ}$ `jump' must occur on angular scales smaller than the beamwidth.338 At ~2/ resolution (about twice tha in Fie. 2)).," At $\sim 2^{\prime}$ resolution (about twice that in Fig. \ref{f-phi349}) ),"339 the ‘canals’ iudeed seeui uuresolved. but the decrease in S/N-ratio precludes couclusious ou even stnaller scales (the original data have 0.8! resolution).," the `canals' indeed seem unresolved, but the decrease in S/N-ratio precludes conclusions on even smaller scales (the original data have $0.8^{\prime}$ resolution)."340 Additional evideuce that the ‘canals are due to beam depolarization is statistical., Additional evidence that the `canals' are due to beam depolarization is statistical.341 We defined *canal-like’ points from the observed values of P. as follows.," We defined `canal-like' points from the observed values of $P$, as follows."342 For cach point in the mosaic we compared the observed value of P with the P-values in pairs of two cliametrically opposed neighbouring (adjacent) poiuts., For each point in the mosaic we compared the observed value of $P$ with the $P$ -values in pairs of two diametrically opposed neighbouring (adjacent) points.343" If the value of P in the ceutral point was less than a certain sinall fraction of the values iu comparison points. the poiut was defined ""eanallike."," If the value of $P$ in the central point was less than a certain small fraction of the values in comparison points, the point was defined `canal-like'."344" This definition ος the visual detection ""aleorithia.", This definition mimics the visual detection `algorithm'.345 Iu the top panel of Fig., In the top panel of Fig.346 3 we show the distribution of the difference between the pol ομως in the two adjaceut points that define the canalblike points. for a P-threshold of30%.," \ref{f-delphi_c} we show the distribution of the difference between the $\phi_{\rm pol}$ 's in the two adjacent points that define the `canal-like' points, for a $P$ -threshold of."347. The Avotistribution peaks at 907. fully consistent with the beam depolarization livpothesis.," The $\Delta\phi_{\rm pol}$ -distribution peaks at $^{\circ}$, fully consistent with the beam depolarization hypothesis."348" This conclusion is reinforced by a comparison with the distribution of Aoy,4; (agn for diametrically opposed adjacent neielibours) of all poiuts for which P? is between 1.0 and 2.0 times lareer than both P-values in the two diametrically opposed neighbouring poiuts. which is shown iu the bottom panel of the same figure."," This conclusion is reinforced by a comparison with the distribution of $\Delta\phi_{\rm pol}$ (again for diametrically opposed adjacent neighbours) of all points for which $P$ is between 1.0 and 2.0 times larger than both $P$ -values in the two diametrically opposed neighbouring points, which is shown in the bottom panel of the same figure."349 Simular ‘canals’ were noted by Uxvaulker et (01999) and Dunucau et ((1998). who also invoked beu depolarization.," Similar `canals' were noted by Uyaniker et (1999) and Duncan et (1998), who also invoked beam depolarization."350 Yet. Fig.," Yet, Fig."351 3 is the first quantitative proof for this explanation., \ref{f-delphi_c} is the first quantitative proof for this explanation.352" Tiwo processes ean cause jnaps in polarization angle oia across the ""canals: a sudden change m RAL across the canals. aud a jump in dufrinsbc oy. of the cuiission incident on the Faraday screei"," Two processes can cause jumps in polarization angle $\phi_{\rm pol}$ across the `canals': a sudden change in RM across the `canals', and a jump in intrinsic $\phi_{\rm pol}$ of the emission incident on the Faraday screen."353 A laree change in intrinsic Opel Muplices a change iu magnetic field direction. aud is therefore quite dificult to understand in view of the absence of structure in total intensity J at the more than evel (see Sect. 2))., A large change in intrinsic $\phi_{\rm pol}$ implies a change in magnetic field direction and is therefore quite difficult to understand in view of the absence of structure in total intensity $I$ at the more than level (see Sect. \ref{s-polint}) ).354 On the other haud. variations iu the RAL of the Faraday screen would seem to be quite natural. if not unavoidable.," On the other hand, variations in the RM of the Faraday screen would seem to be quite natural, if not unavoidable."355" Discoutinuities iu RAL aust plav an important rolle in producing the ""canals; because the ""canals; although siuuilar in adjacent frequency bands. generally do not occur in all bands. ancl certaimly are not ideutical iu the different bauds (see Fig. 2))."," Discontinuities in RM must play an important rôlle in producing the `canals', because the `canals', although similar in adjacent frequency bands, generally do not occur in all bands, and certainly are not identical in the different bands (see Fig. \ref{f-phi349}) )."356" This indicates that the jumps IM Oy, ave mainly due to changes in RM.", This indicates that the jumps in $\phi_{\rm pol}$ are mainly due to changes in RM.357" Towever. the question is if the Jumps in ©.) are indeed accompanied by juups in RM so that Aou,=907 is xoduced at the frequency where the ‘canal’ is best visible."," However, the question is if the jumps in $\phi_{\rm pol}$ are indeed accompanied by jumps in RM so that $\Delta\phi_{\rm pol} = 90^{\circ}$ is produced at the frequency where the `canal' is best visible."358 Tn principle. the determination of RAL ouly involves a sluuple linear fit of the polarization angles in the five ποιος bauds (at 311. 319. 355. 360 and 375 MIIZ) A2. mt iu practice several complications may arise.," In principle, the determination of RM only involves a simple linear fit of the polarization angles in the five frequency bands (at 341, 349, 355, 360 and 375 MHz) $\lambda^2$, but in practice several complications may arise."359" First. he observed values of oo, may be biased due to inagiug effects (like off-sets} in the Stokes Q- and Canaps from which oo. is derived WWicringa ot 11993)."," First, the observed values of $\phi_{\rm360pol}$ may be biased due to imaging effects (like off-sets) in the Stokes $Q$ - and $U$ -maps from which $\phi_{\rm pol}$ is derived Wieringa et 1993)."361 Our data indicate that. in the maps of this region of sky. such," Our data indicate that, in the maps of this region of sky, such"362of open clusters in the evolution of the stellar population of the Galactic disc.,of open clusters in the evolution of the stellar population of the Galactic disc.363 The paper has the following structure., The paper has the following structure.364 In Sect., In Sect.365 ?? we briefly describe the input data and give the definitions used throughout the paper., \ref{sec:dadef} we briefly describe the input data and give the definitions used throughout the paper.366 Sect., Sect.367 ?? is devoted to the construction of cluster the luminosity and mass functions., \ref{sec:lmf} is devoted to the construction of cluster the luminosity and mass functions.368 In Sect., In Sect.369 4 we examine how cluster mass and luminosity functions evolve with time and build the cluster initial mass function., \ref{sec:evol} we examine how cluster mass and luminosity functions evolve with time and build the cluster initial mass function.370 In Sects., In Sects.371 ?? and ?? we discuss and summarise the results., \ref{sec:discuss} and \ref{sec:concl} we discuss and summarise the results.372 In this paper we use the results of the previous work based on the study of 650 nearby open clusters identified in the ASCC-2.5., In this paper we use the results of the previous work based on the study of 650 nearby open clusters identified in the ASCC-2.5.373 The sample includes 520 known objects and 130 open clusters detected within our project (?).., The sample includes 520 known objects and 130 open clusters detected within our project \citep[]{newclu}.374" For each star projected on a cluster area, a membership probability was determined in an iterative process that takes spatial, photometric, and proper motion distributions of stars into account within the corresponding area on the sky (?).."," For each star projected on a cluster area, a membership probability was determined in an iterative process that takes spatial, photometric, and proper motion distributions of stars into account within the corresponding area on the sky \citep[]{starcat}."375" At the end of the iterations for each cluster, we obtained new coordinates of the cluster centre, the cluster size, the mean proper motion, the distance from the Sun, reddening, and age (??).."," At the end of the iterations for each cluster, we obtained new coordinates of the cluster centre, the cluster size, the mean proper motion, the distance from the Sun, reddening, and age \citep[][]{clucat,newclu}."376" These parameters were determined with data on the most probable members, i.e., stars having both kinematic and photometric membership probabilities higher than61%."," These parameters were determined with data on the most probable members, i.e., stars having both kinematic and photometric membership probabilities higher than."377. The results are included in the Catalogue of Open Cluster Data (COCD) and its Extension (??)..," The results are included in the Catalogue of Open Cluster Data (COCD) and its Extension \citep{clucat,newclu}."378 In ? we published masses of 236 clusters of our sample estimated from tidal radii determined with a three-parameter fit of King's profiles to the observed density distributions of cluster members., In \citet{clumart} we published masses of 236 clusters of our sample estimated from tidal radii determined with a three-parameter fit of King's profiles to the observed density distributions of cluster members.379" To obtain mass estimates for all clusters, we used these data to establish a relation between tidal radius and the observed semi-major axis of the apparent distribution of cluster members on the sky."," To obtain mass estimates for all clusters, we used these data to establish a relation between tidal radius and the observed semi-major axis of the apparent distribution of cluster members on the sky."380 The resulting tidal masses (i.e. calibrated tidal masses) for 650 clusters have been published in ?.., The resulting tidal masses (i.e. calibrated tidal masses) for 650 clusters have been published in \citet{clumart1}.381" Though calibrated tidal masses of individual clusters are less accurate, their distribution fits well the distribution of the tidal masses based on the direct fitting of the King model."," Though calibrated tidal masses of individual clusters are less accurate, their distribution fits well the distribution of the tidal masses based on the direct fitting of the King model."382" Moreover, because the parameters are homogeneous and more numerous, the calibrated masses are better suited to statistical investigations."," Moreover, because the parameters are homogeneous and more numerous, the calibrated masses are better suited to statistical investigations."383" In the following, we use the data on the calibrated tidal masses to study the luminosity and mass functions of the Galactic open clusters."," In the following, we use the data on the calibrated tidal masses to study the luminosity and mass functions of the Galactic open clusters."384" Throughout the paper, masses are given in the units of solar masses."," Throughout the paper, masses are given in the units of solar masses."385" Our list of 650 Galactic open clusters identified in the ASCC-2.5 includes 9 cluster-like associations (e.g. Vel OB2, Sco OB4), which require further commentary."," Our list of 650 Galactic open clusters identified in the ASCC-2.5 includes 9 cluster-like associations (e.g. Vel OB2, Sco OB4), which require further commentary."386" These are objects, which are classified both as clusters and associations in the literature (e.g. a Per = Per OB3)."," These are objects, which are classified both as clusters and associations in the literature (e.g. $\alpha$ Per = Per OB3)."387" They are included in our sample because their observed properties (spatial structure, internal kinematics, and stellar content) make them indistinguishable from regular open clusters, although they have larger sizes and higher masses."," They are included in our sample because their observed properties (spatial structure, internal kinematics, and stellar content) make them indistinguishable from regular open clusters, although they have larger sizes and higher masses."388 Hereafter we refer to them as cluster-like associations (seealso?).., Hereafter we refer to them as cluster-like associations \citep[see also ][]{newrv07}.389" We consider two young clusters, NGC 869 (also known as h Per) and NGC 884 (y Per), as a single entity."," We consider two young clusters, NGC 869 (also known as $h$ Per) and NGC 884 $\chi$ Per), as a single entity."390" They overlap in the projection on the sky, share a huge corona, and have a large number (more than 50%)) of members in common, making an accurate determination of their individual tidal radii and masses rather difficult."," They overlap in the projection on the sky, share a huge corona, and have a large number (more than ) of members in common, making an accurate determination of their individual tidal radii and masses rather difficult."391" Also, we exclude the cluster Mamajek 1 (7 Chamaeleontis) for which we were able to identify only three members in theASCC-2."," Also, we exclude the cluster Mamajek 1 $\eta$ Chamaeleontis) for which we were able to identify only three members in the."392"5.. According to Fig. 2,,"," According to Fig. \ref{fig:agemass},"393 this cluster can be omitted without consequences for the results., this cluster can be omitted without consequences for the results.394" Therefore, our final cluster sample includes 648 entities."," Therefore, our final cluster sample includes 648 entities."395" As a measure of cluster brightness needed for the construction of the luminosity function, we take the integrated magnitudes of the clusters."," As a measure of cluster brightness needed for the construction of the luminosity function, we take the integrated magnitudes of the clusters."396" Doing so, we can directly compare our results with findings based on observations of open clusters in other galaxies."," Doing so, we can directly compare our results with findings based on observations of open clusters in other galaxies."397" We define apparent and absolute integrated magnitudes Jy and Jy, as and where N; and V; are the number and the apparent magnitudes of the most probable cluster members, respectively, and (V—My) is the apparent distance modulus available in the for each cluster."," We define apparent and absolute integrated magnitudes $I_V$ and $I_{M_V}$ as and where $N_1$ and $V_i$ are the number and the apparent magnitudes of the most probable cluster members, respectively, and $(V-M_V)$ is the apparent distance modulus available in the for each cluster."398" For “unseen” stars, Aly is a magnitude correction, introduced to make the computed /y and Iy, independent of the range of stellar magnitudes actually Observed in a given cluster."," For “unseen” stars, $\Delta I_V$ is a magnitude correction, introduced to make the computed $I_V$ and $I_{M_V}$ independent of the range of stellar magnitudes actually observed in a given cluster."399"It is naturally expected then, that stellar clusters are affected by DM halos, since their basic constituents, namely stars, are themselves affected.","It is naturally expected then, that stellar clusters are affected by DM halos, since their basic constituents, namely stars, are themselves affected."400 The main reason is the fact that stars with lower masses evolve slower in dense DM halos., The main reason is the fact that stars with lower masses evolve slower in dense DM halos.401" This effect is not noticeable for young clusters since in these clusters low-mass stars are still in the MS and the more massive ones, which are evolving through the RGB, are not affected by the presence of DM."," This effect is not noticeable for young clusters since in these clusters low-mass stars are still in the MS and the more massive ones, which are evolving through the RGB, are not affected by the presence of DM."402" However, in old clusters the RGB may be populated by stars that evolved slower, consequently making the cluster look younger than its real age."," However, in old clusters the RGB may be populated by stars that evolved slower, consequently making the cluster look younger than its real age."403" Moreover, the fact that low-mass stars within dense DM halos follow brighter paths in the RGB than classical stars contributes to amplify this effect."," Moreover, the fact that low-mass stars within dense DM halos follow brighter paths in the RGB than classical stars contributes to amplify this effect."404" In order to distinctly illustrate the younger appearance of a cluster when embedded in a dense DM halo, we computed the isochrones (the track drawn by the positions in the H-R diagram of all stars with different masses at a given of stellar clusters in different situations."," In order to distinctly illustrate the younger appearance of a cluster when embedded in a dense DM halo, we computed the isochrones (the track drawn by the positions in the H-R diagram of all stars with different masses at a given age) of stellar clusters in different situations."405 Figure 2 age)shows the isochrones we obtained for a cluster evolving in a halo of DM with a density px=10° GeV cm lines) together with those obtained without? the (continuousinfluence of DM (dashed lines)., Figure \ref{fig-isoe09} shows the isochrones we obtained for a cluster evolving in a halo of DM with a density $\rho_{\chi}=10^9\;$ GeV $^{-3}$ (continuous lines) together with those obtained without the influence of DM (dashed lines).406" When the isochrones of >1000 Myr in both situations are compared, we see that indeed the cluster within a dense DM halo looks younger, with a brighter and hotter turn-off point and a brighter RGB."," When the isochrones of $\geq 1000\;$ Myr in both situations are compared, we see that indeed the cluster within a dense DM halo looks younger, with a brighter and hotter turn-off point and a brighter RGB."407" In this case the turn-off and RGB are populated by more massive stars than in the classical scenario, because they took longer to burn out their hydrogen core and to leave the MS."," In this case the turn-off and RGB are populated by more massive stars than in the classical scenario, because they took longer to burn out their hydrogen core and to leave the MS."408 It is almost impossible to distinguish both clusters at ages <500 Myr., It is almost impossible to distinguish both clusters at ages $\leq 500\;$ Myr.409" When even higher DM densities are considered (or, equivalently, larger WIMP-on-nucleon scattering cross sections) the characteristics of the cluster change dramatically."," When even higher DM densities are considered (or, equivalently, larger WIMP-on-nucleon scattering cross sections) the characteristics of the cluster change dramatically."410" In addition to the previously described effect will now be visible for younger clusters, because at (whichhigher DM densities more massive stars will be affected), another strong signature of the presence of DM in the halo arises when looking at the position of stars with lower masses."," In addition to the previously described effect (which will now be visible for younger clusters, because at higher DM densities more massive stars will be affected), another strong signature of the presence of DM in the halo arises when looking at the position of stars with lower masses."411" These stars, which are mostly fueled by the energy from DM annihilation, go back in the Hyashi track and reach positions in the H-R diagram which were normally occupied only by forming stars in their way to the MS."," These stars, which are mostly fueled by the energy from DM annihilation, go back in the Hyashi track and reach positions in the H-R diagram which were normally occupied only by forming stars in their way to the MS."412" Consequently, the bottom of the isochrones, corresponding to the lower mass stars, rises to higher luminosities, giving the cluster a very characteristic appearance."," Consequently, the bottom of the isochrones, corresponding to the lower mass stars, rises to higher luminosities, giving the cluster a very characteristic appearance."413 This peculiar signature is a strong indication of the presence of high concentrations of DM in a stellar cluster., This peculiar signature is a strong indication of the presence of high concentrations of DM in a stellar cluster.414" This strong signature is illustrated in Figure 3,, where the isochrones of a stellar cluster surrounded by a halo of DM with a density p,=1010 GeV cm~® are plotted."," This strong signature is illustrated in Figure \ref{fig-isoe10}, where the isochrones of a stellar cluster surrounded by a halo of DM with a density $\rho_{\chi}=10^{10}\;$ GeV $^{-3}$ are plotted."415 The main characteristic signature of the presence of DM is the fact that the bottom of all isochrones is more than 3 times brighter than the classical isochrones., The main characteristic signature of the presence of DM is the fact that the bottom of all isochrones is more than 3 times brighter than the classical isochrones.416" In addition, the effect of a brighter and hotter turn-off point is now more pronounced and appreciable in clusters as young as 250 Myr."," In addition, the effect of a brighter and hotter turn-off point is now more pronounced and appreciable in clusters as young as 250 Myr."417 We have also considered the hypothetical scenario in which DM is formed by the low-mass WIMPs invoked to reconcile the results of DAMA with the negative results of other direct detection experiments 2009).., We have also considered the hypothetical scenario in which DM is formed by the low-mass WIMPs invoked to reconcile the results of DAMA with the negative results of other direct detection experiments \citep{art-Savageetal2009JCAP}.418" As shown in Figure 4,, if such WIMPs (Savageform most of the DM then the DM density needed to have signatures"," As shown in Figure \ref{fig-mx8}, if such WIMPs form most of the DM then the DM density needed to have signatures"419from Eq.CÀ5)).,from \ref{A5}) ).420 Particularly we are interesting in asvinptotic No»0. which should reflect directly the svuunetry of the foreground signal.," Particularly we are interesting in asymptotic $N\rightarrow 0$, which should reflect directly the symmetry of the foreground signal."421 Simple algebra allows us to represeut No function in the following form Taking into account that area S is located close to the 0=7/2. let us discuss the properties of N-fuuction at the lint (25/251. using asvinptotic of the Legeudre polwnomials.," Simple algebra allows us to represent $N$ function in the following form Taking into account that area $S$ is located close to the $\theta=\pi/2$, let us discuss the properties of N-function at the limit $\l\pi/2\gg1$, using asymptotic of the Legendre polynomials."422 After simple algebra we obtain where From Eq.(As)) oue cau see tle syumuctry of theLegendre polynomials which manifest themself trough οὐκ aud zu0;A modes., After simple algebra we obtain where From \ref{A8}) ) one can see the symmetry of theLegendre polynomials which manifest themselfs trough $\cos\theta_j\Delta$ and $\sin\theta_j\Delta$ modes.423 If0;=7/2. then depeudiug on A we will lave Thus. choosing A=ln.1.2... mode we take the correspouding properties of the Legeudre polvuocmials into consideration.," If$\theta_j=\pi/2$, then depending on $\Delta$ we will have Thus, choosing $\Delta=4n, n=1,2...$ mode we take the corresponding properties of the Legendre polynomials into consideration."424 However. as one can see from Eq.(A9)) A=lao=1.2... periodicity of the Galactic image is not exact.," However, as one can see from \ref{A9}) ) $\Delta=4n, n=1,2...$ periodicity of the Galactic image is not exact."425 Iu reality we have pixels which coutaiut galaxy signal having theta; coordinate close to 7/2. but not exactly equivaleut to m/2.," In reality we have pixels which containt galaxy signal having $theta_j$ coordinate close to $\pi/2$, but not exactly equivalent to $\pi/2$ ."426 Let us introduce a now variable 0;=7/2.0;.6;<1 characterized the deviation of the j-th pixel location foin the theta=7/2 planc.," Let us introduce a new variable $\delta_j= \pi/2-\theta_j, \delta_j\ll 1$ characterized the deviation of the $j$ -th pixel location from the $theta=\pi/2$ plane."427 From Eq.(As8)) one can find that for A=lv.n1.2... cosdAC9CONLid;cd(lid;}22.," From \ref{A8}) ) one can find that for $\Delta=4n, n=1,2...$ $\cos\theta_j\Delta\simeq \cos4n\delta_j\simeq 1-(4n\delta_j)2/2$."428 Thus. if pixel j coutaimt the igual from the galactie foreground. the deviation from the center of the ealactic plane should be siull enough;<a/le.," Thus, if pixel $j$ containt the signal from the galactic foreground, the deviation from the center of the galactic plane should be small $\delta_j\ll \pi/4n$."429 It is clear that this condition does not uccessarily correspond to the properties of the Calactie mage. which is clearly secu from the Ik. ἵνα aud Q baud signals.," It is clear that this condition does not necessarily correspond to the properties of the Galactic image, which is clearly seen from the K, Ka and Q band signals."430" Taking the aboveanenutioned. properties of GIA,(0;0,) faction. we represent the asviuptotic of this function at the limit 6;<z/bi. which is applicable for analysis of the Galactic signal at V and W bauds."," Taking the above-mentioned properties of $G^m_{\l,\Delta=4n}(\theta_j,\theta_k)$ function, we represent the asymptotic of this function at the limit $\delta_j\ll \pi/4n$, which is applicable for analysis of the Galactic signal at V and W bands."431 Thus. combining Eq.CÀT)) aud ALO)). we obtain One may think that choose of described above la-mode of Legendre polvionmials automatically euarautec cancellation of the brightest part of the signal from the map without auv restriction on sviunietrv aud amplitude of the foreeround.," Thus, combining \ref{A7}) ) and \ref{A10}) ), we obtain One may think that choose of described above $4n$ -mode of Legendre polynomials automatically guarantee cancellation of the brightest part of the signal from the map without any restriction on symmetry and amplitude of the foreground."432 To show that svuunetry of the Galactic signal is important. let us discuss a few particular cases; which illuminate this problem more clearly.," To show that symmetry of the Galactic signal is important, let us discuss a few particular cases, which illuminate this problem more clearly."433 Firstly. let take a look at ealactic ceuter (GC). which is one of the brightest sources of the signal.," Firstly, let take a look at galactic center (GC), which is one of the brightest sources of the signal."434 For the GC corresponding auplitiudes 7; are localized per pixels. for which o;~0 in the galactic svstem of coordinate.," For the GC corresponding amplitudes $T_j$ are localized per pixels, for which $\phi_j\simeq 0$ in the galactic system of coordinate."435 From ALL))) one can see that for GC the πιοοι Vis equivalent to zero., From \ref{A11}) ) one can see that for GC the function $N$ is equivalent to zero.436 More accurately. talking iuto account that image of the (GC has characteristic sizes 00=ào~EWIINL where FWIIM is the Full Width of Half Maxiuuni of the beam. in Eq.(AL!—) additionally to 6;«1 parameter we got small paraucter mFWOAL<1.," More accurately, taking into account that image of the GC has characteristic sizes $\delta\theta=\delta\phi\sim {\rm FWHM}$, where FWHM is the Full Width of Half Maximum of the beam, in \ref{A11}) ) additionally to $\delta_j\ll 1$ parameter we get small parameter $m{\rm FWHM}\ll 1$."437 Secondly. let discuss the model of two bright point like sources. located svuuuctrically relatively to the GC.," Secondly, let discuss the model of two bright point like sources, located symmetrically relatively to the GC."438 Let assuues that for that point sources Zi4T5. but ο04=s.," Let assumes that for that point sources $T_1\neq T_2$, but $\phi_2-\phi_1=\pi$."439 Όπου again. from Eq.(AL1)) we ect No=0 for all ij and these poiut sources will be automatically removed by the di estimator even if they lave 944s.," Once again, from \ref{A11}) ) we get $N=0$ for all $m$ and these point sources will be automatically removed by the $d^{4n}_{\l,m}$ estimator even if they have $\delta_1\neq\delta_2$."440 Another possibility related to the svimmetry of the Calactic mage in 0 direction., Another possibility related to the symmetry of the Galactic image in $\theta$ direction.441 We would like to remind. that Eq.CÀTI)) was obtained under approsimation (0; <1. where 0;=πο 0;.," We would like to remind, that \ref{A11}) ) was obtained under approximation $\l\theta_j\ll 1$ , where $\theta_j=\pi/2-\delta_j$ ."442 This means. that (6; ein be asbig enough ((4)z 1). as sunall (68;< 1) as well.," This means, that $\l\delta_j$ can be asbig enough $\l \delta_j \gg 1$ ), as small $\l\delta_j\ll1$ ) as well."443" For (6)<1 from Eq.(AL0)) we obtain As one can see from Eq.CÀA12)) the bright sources located on the same ϐ coordinates (0;= ó,) does not contribute to di"" estimator."," For $\l\delta_j\ll 1$ from \ref{A10}) ) we obtain As one can see from \ref{A12}) ) the bright sources located on the same $\theta$ coordinates $\delta_j=\delta_k$ ) does not contribute to $d^{4n}_{\l,m}$ estimator."444The advent of 102m class telescopes in the last decade has provided us with an unprecedented view of the high-redshift universe.,The advent of 10-m class telescopes in the last decade has provided us with an unprecedented view of the high-redshift universe.445 In particular. LILRES. the high-resolution echelle spectrograph on the Ixeck E telescope in Hawaii. has mace it possible to obtain observations with resolution as high as ~6+.," In particular, HIRES, the high-resolution echelle spectrograph on the Keck I telescope in Hawaii, has made it possible to obtain observations with resolution as high as $\sim 6$."446 As a result. a number of detailed studies of absorption lines in the spectra of background. QSOs have been carried out (for reviews see Rauch 1998: Hamann Ferlanel 1999).," As a result, a number of detailed studies of absorption lines in the spectra of background QSOs have been carried out (for reviews see Rauch 1998; Hamann Ferland 1999)."447 Llowever. this improvement inspecfrad resolution cannot be casily accompanied by an improvement inspatial resolution: 2-dimoensional information about the absorbers can normally only be obtained by comparing separate spectra of multiple quasar images. which are duc to cither gravitationally lensed. single QSOs or actual QSO pairs which have a small angular separation on the plane of the sky.," However, this improvement in resolution cannot be easily accompanied by an improvement in resolution: 2-dimensional information about the absorbers can normally only be obtained by comparing separate spectra of multiple quasar images, which are due to either gravitationally lensed single QSOs or actual QSO pairs which have a small angular separation on the plane of the sky."448 Because lensed QSOs have image separations up o a [ew areseconds. they are useful for probing scales =1090 kpe.," Because lensed QSOs have image separations up to a few arcseconds, they are useful for probing scales $\approxlt 100$ kpc."449 QSO pairs have image separations of at least a few areminutes. thus allowing one to probe scales of a ow hundred kpe.," QSO pairs have image separations of at least a few arcminutes, thus allowing one to probe scales of a few hundred kpc."450 From coincident absorption in more than one line of sight (LOS). minimum size (coherence length) estimates can be obtained. immediately.," From coincident absorption in more than one line of sight (LOS), minimum size (coherence length) estimates can be obtained immediately."451 Estimates of most xobable sizes require a statistical “hit and aapproach (e.g. MeCGill 1990)., Estimates of most probable sizes require a statistical hit and approach (e.g. McGill 1990).452 Such studies have lec to size estimates for aabsorbers ranging [rom a lew kpe (Κος ct al., Such studies have led to size estimates for absorbers ranging from a few kpc (Foltz et al.453 1984: AMeCGuill 1990: Bechtolel Yee 1995) to a few hundred kpc (Dinshaw et al., 1984; McGuill 1990; Bechtold Yee 1995) to a few hundred kpc (Dinshaw et al.454 1995: Smette et al., 1995; Smette et al.455 1995: Dinshaw et al., 1995; Dinshaw et al.456 1998) to ~1 Alpe (Fang et al., 1998) to $\sim 1$ Mpc (Fang et al.457 1996: Dinshaw et al., 1996; Dinshaw et al.458 1997: Young et al., 1997; Young et al.459 2001)., 2001).460 Such Large ssugeest that the typical. low column density 1014.5 τὸ) [forest line. does not originate in cliscrete galactic or protogalactic clouds.," Such large suggest that the typical, low column density $\approxlt 10^{14.5}$ ) forest line does not originate in discrete galactic or protogalactic clouds."461 The emerging picture is consistent with results from hyvdrodynamical numerical simulations in a Cold Dark Matter (CDM) structure formation scenario according to which the [forest is a manifestation of a space filling photoionized, The emerging picture is consistent with results from hydrodynamical numerical simulations in a Cold Dark Matter (CDM) structure formation scenario according to which the forest is a manifestation of a space filling photoionized462"115 + 20 for the D; and Ds lines, respectively.","115 $\pm$ 20 for the $_1$ and $_2$ lines, respectively."463 The calibration of ? gives E(B— V)= 0.03 + 0.01., The calibration of \cite{MunariZwitter1997} gives $E(B-V)$ = 0.03 $\pm$ 0.01.464 'The dereddened spectral energy distribution was fitted using solar-composition (7) model fluxes., The dereddened spectral energy distribution was fitted using solar-composition \citep{Kurucz1993} model fluxes.465 The model fluxes were convolved with photometric filter response functions., The model fluxes were convolved with photometric filter response functions.466 A weighted Levenberg-Marquardt non-linear least-squares fitting procedure was used to find the solution that mimimized the difference between the observed and model fluxes., A weighted Levenberg-Marquardt non-linear least-squares fitting procedure was used to find the solution that mimimized the difference between the observed and model fluxes.467" Since the surface gravity is poorly constrained by our spectral energy distribution, fits were performed for logg—4.5 and logg—2 to assess the uncertainty due to unconstrained logg."," Since the surface gravity is poorly constrained by our spectral energy distribution, fits were performed for $\log g = 4.5$ and $\log g = 2$ to assess the uncertainty due to unconstrained $\log g$."468 A final value of Teg=7140-310 K was found., A final value of $T_{\rm eff} = 7140 \pm 310$ K was found.469 The uncertainties in Tig includes the formal least-squares error and that from the uncertainties in E(B—V) and logg., The uncertainties in $T_{\rm eff}$ includes the formal least-squares error and that from the uncertainties in $E(B-V)$ and $\log g$.470 We note here that in the next section values with considerably smaller uncertainties are determined from high-dispersion spectroscopy., We note here that in the next section values with considerably smaller uncertainties are determined from high-dispersion spectroscopy.471 In order to classify the star with higher precision and, In order to classify the star with higher precision and472where dft is the time elapsed since 2155568.05 BID. ic. the time at which the background increases sharply (Fie.,"where $dt$ is the time elapsed since 2455568.05 BJD, i.e. the time at which the background increases sharply (Fig."473 1 2)., 1 2).474 The ayy and ay terms were ouly applied for df>0., The $a_{11}$ and $a_{12}$ terms were only applied for $dt > 0$.475" The six first terms of this equation correspond to the ""Upixebl-phase model (Eq.", The six first terms of this equation correspond to the “pixel-phase” model (Eq.476 1)., 1).477 The purpose of the linear enu in df is to model a possible smooth variation of he stellar brightness., The purpose of the linear term in $dt$ is to model a possible smooth variation of the stellar brightness.478 The other terms result from our extensive analysis of our cutive set of wariSpitser data (Cullon et al.," The other terms result from our extensive analysis of our entire set of warm data (Gillon et al.,"479 in prep.), in prep.)480 aud of available calibration data hat lead us to conclude to a low-aiplitude periodic variability of the effective gan of the detector. its typical xeriod Wing between 30 and GO minutes and its amplitude cine iu average of a few dozeus of ppm.," and of available calibration data that lead us to conclude to a low-amplitude periodic variability of the effective gain of the detector, its typical period lying between 30 and 60 minutes and its amplitude being in average of a few dozens of ppm."481 Considering he challenging photometric precision required by our ΟΤΕ. it is very muportant to take it iuto account. justifvine the sine teziu in Eq.," Considering the challenging photometric precision required by our program, it is very important to take it into account, justifying the sine term in Eq."482 2., 2.483 We are currently working with theSpitzer cuginecring team to find the origin of his periodic variation., We are currently working with the engineering team to find the origin of this periodic variation.484" We also uotice that à ""background explosion” such as the one affecting the last part of our data is correlated to a sharp increase of the effective eaiu of the detector that is very well modeled by 1ὸ last two terms of Eq.", We also notice that a “background explosion” such as the one affecting the last part of our data is correlated to a sharp increase of the effective gain of the detector that is very well modeled by the last two terms of Eq.485 2., 2.486 The AICAIC uses the whol dataset to siunultaueouslv fit for the transit model aud jo baseline function preseuted in Eq., The MCMC uses the whole dataset to simultaneously fit for the transit model and the baseline function presented in Eq.487 2., 2.488" The following parameters were jump πι our analysis:+ the planot/starH area ratioB P(R,/R.j-. the trausit width (from first to last contact) WV. the impact paralucter b=0cos//RÉ.. aud the time of ninimiuaua light Jy."," The following parameters were jump in our analysis: the planet/star area ratio $(R_p /R_s )^2$, the transit width (from first to last contact) $W$, the impact parameter $b = a \cos{i}/R_\ast$, and the time of minimum light $T_0$."489 We asstuned a uniform prior distribution for these juup parameters. but we imposed a Cassia prior for the stellar radius 2 based. on %=0.913001601. (Table 1).," We assumed a uniform prior distribution for these jump parameters, but we imposed a Gaussian prior for the stellar radius $R_\ast$ based on $R_\ast = 0.943 \pm 0.010 R_\odot$ (Table 1)."490 We assumed a quadratic lub-darkening lav with coefficients 44=0.0706 and uy=0.1171., We assumed a quadratic limb-darkening law with coefficients $u_1=0.0706$ and $u_2=0.1471$.491 These values were drawn roni the theoretical tables of Claret Blocmen (2011) for the TRAC 15 fan bandpass aud for T= 5250 FK. =L5 and [Fe/TI]|210.3.," These values were drawn from the theoretical tables of Claret Bloemen (2011) for the IRAC 4.5 $\mu$ m bandpass and for = 5250 K, =4.5 and [Fe/H]=+0.3."492 The 12 cocfiicicuts of the baseline models (Eq., The 12 coefficients of the baseline models (Eq.493 2) were determined by least-square mininization at each steps of the Markov chains (see (10 and references therein for details)., 2) were determined by least-square minimization at each steps of the Markov chains (see G10 and references therein for details).494" The correlated) noise present im the LC was taken iuto account as describe by CLO. ie.. a scalime factor jJ, was determined from the standard deviation of the binned aud unbiunued residuals of a preliminary MCMC analysis. aud it was applied to the error bars."," The correlated noise present in the LC was taken into account as described by G10, i.e., a scaling factor $\beta_r$ was determined from the standard deviation of the binned and unbinned residuals of a preliminary MCMC analysis, and it was applied to the error bars."495 Several biuniug iutervals ranging from 10 to 90 nüuutes were tried in proelinimiuv short Markov Chains. aud the maximal value for 3). 1.35. was used m our analysis.," Several binning intervals ranging from 10 to 90 minutes were tried in preliminary short Markov Chains, and the maximal value for $\beta_r$, 1.35, was used in our analysis."496 We performed two new MCMC analyses. one with a transit of CCnceec. and one without.," We performed two new MCMC analyses, one with a transit of e, and one without."497 Figure 1| shows the resulting best-fit transit model and its residuals., Figure 4 shows the resulting best-fit transit model and its residuals.498 The odds ratio (Eq., The odds ratio (Eq.499 2 alone) es (Eq., 2 alone) $vs$ (Eq.500 2 transl)ds ~107 in favor of the transit model., 2 + transit) is $\sim 10^8$ in favor of the transit model.501 The transit of COncoe is thus firmly detected., The transit of e is thus firmly detected.502 The period of the siuusoid (040) derived roni the MCAIC is 51 minutes. we decoupled from the trausit duration (96 miuutes) aud significantly longer than the pixelphase timescale (35 iumnutes).," The period of the sinusoid $a_{10}$ ) derived from the MCMC is 51 minutes, well decoupled from the transit duration (96 minutes) and significantly longer than the pixel-phase timescale (35 minutes)."503 Its auplitude is 115427 ppii., Its amplitude is $\pm$ 27 ppm.504 We show in Fie., We show in Fig.505 5 he differeut contributions of the spatially aud ine-depeudeut ternis of Eq., 5 the different contributions of the spatially- and time-dependent terms of Eq.506 2., 2.507 This shows that the time-depeudent teris are well decoupled from the transit pattern., This shows that the time-dependent terms are well decoupled from the transit pattern.508 Table 2 preseuts the resultingC» transit aud plivsical parameters aud Lo error lanits derived for ec. We also conducted. a residual permutation bootstrap analysis. known as the praver bead method (Callon et al.," Table 2 presents the resulting transit and physical parameters and $\sigma$ error limits derived for e. We also conducted a residual permutation bootstrap analysis, known as the prayer bead method (Gillon et al."509 2006) to obtain au additional estimation of the residual correlated noise., 2006) to obtain an additional estimation of the residual correlated noise.510 We used for this purpose the liehteurve corrected from the svsteimatic effects described in Eq., We used for this purpose the lightcurve corrected from the systematic effects described in Eq.511 2., 2.512 The resulting parameters are in excellent aereement with the ones derived from the AICAIC analysis (Table 2). while their error bars are significantly smaller.," The resulting parameters are in excellent agreement with the ones derived from the MCMC analysis (Table 2), while their error bars are significantly smaller."513 This result indicates that the eror budget is dominated bv the nucertaintics ou the parameters of the complex baseline model. aud not by the residual correlated noise.," This result indicates that the error budget is dominated by the uncertainties on the parameters of the complex baseline model, and not by the residual correlated noise."514 To test the robustness of our transit detection amd of the resulting transit parameters. we performed ~ 10 additional ATCAIC analyses as described above. each of them assuming a different set of time-dependent ternis presented in Eq.," To test the robustness of our transit detection and of the resulting transit parameters, we performed $\sim$ 10 additional MCMC analyses as described above, each of them assuming a different set of time-dependent terms presented in Eq."515 2., 2.516 We used a binned lightcurve (per 30s) for the purpose of this comparison to speed up the analysis., We used a binned lightcurve (per 30s) for the purpose of this comparison to speed up the analysis.517 Table 3Pa preseuts the baseline model. derived," Table 3 presents the baseline model, derived"518where the subscripts 0 and 1 denote (the equilibrium solutions and the perturbations. respectively.,"where the subscripts $0$ and $1$ denote the equilibrium solutions and the perturbations, respectively."519 Using equations (25)) and (26)). we linearize equations (14)) aud (15)) and obtain and We consider plane-wave solutions xexp(/pz—Aet) Dor ó4 and e. where w is (he angular frequency and p is the wave number.," Using equations \ref{eq:phi1}) ) and \ref{eq:psi1}) ), we linearize equations \ref{eq:eqmot1}) ) and \ref{eq:eqmot2}) ) and obtain and We consider plane-wave solutions $\propto \exp(ipz-i\omega t)$ for $\phi_1$ and $\psi_1$, where $\omega$ is the angular frequency and $p$ is the wave number."520 Then. we derive the dispersion relation for oscillations. If we adopt such solutions as (y=—6o for equilibrium. from equation (29)) we obtain the dispersion relation simplified as Next. we examine the property of oscillations represented by the dispersion relation (30)).," Then, we derive the dispersion relation for oscillations, If we adopt such solutions as $\psi_0=- \phi_0$ for equilibrium, from equation \ref{eq:disp}) ) we obtain the dispersion relation simplified as Next, we examine the property of oscillations represented by the dispersion relation \ref{eq:disp2}) )."521 The first term in the dispersion relation (20)) can be derived when in equations (27)) and (28)) we include only the second and third. terms ( the Magnus ancl (tension terms) ignoring the fist term (the pinning term)., The first term in the dispersion relation \ref{eq:disp2}) ) can be derived when in equations \ref{eq:pert1}) ) and \ref{eq:pert2}) ) we include only the second and third terms ( the Magnus and tension terms) ignoring the first term (the pinning term).522 Hence. the first term in the dispersion relation (30)) expresses the Ixelvin wave. whichis circularly polarized ancl propagates helically along a vortex line (Sonin1957).," Hence, the first term in the dispersion relation \ref{eq:disp2}) ) expresses the Kelvin wave, whichis circularly polarized and propagates helically along a vortex line \citep{son87}."523. The second term in the dispersion relation (30)) can be derivel when in equations (27)) aud (28)) we include only (he first and second terms (the pinning and Magnus terms) neglecting the third term (the tension term)., The second term in the dispersion relation \ref{eq:disp2}) ) can be derived when in equations \ref{eq:pert1}) ) and \ref{eq:pert2}) ) we include only the first and second terms (the pinning and Magnus terms) neglecting the third term (the tension term).524 We find that the second term in (he dispersion relation expresses the circular motion of a vortex line as à whole around the equilibrium position., We find that the second term in the dispersion relation expresses the circular motion of a vortex line as a whole around the equilibrium position.525 In general the dispersion relation (230)) represents the combined oscillations of Kelvin and rotational modes., In general the dispersion relation \ref{eq:disp2}) ) represents the combined oscillations of Kelvin and rotational modes.526 Figure 6 illustrates the dispersion relation given by equation (30)) for different oy., Figure \ref{disp} illustrates the dispersion relation given by equation \ref{eq:disp2}) ) for different $\phi_0$.527 The angular frequeney of allowed oscillations is plotted as a funetion of wave number., The angular frequency of allowed oscillations is plotted as a function of wave number.528 At smaller wave numbers (he angular frequency. of oscillations is determined mainly by the rotational mode. while at larger wave numbers maiulv by the Ixelvin mode.," At smaller wave numbers the angular frequency of oscillations is determined mainly by the rotational mode, while at larger wave numbers mainly by the Kelvin mode."529 We should note that there exists (he minimum frequency for the oscillations especially when ὦμ is close to 2»., We should note that there exists the minimum frequency for the oscillations especially when $\phi_0$ is close to $2n\pi $ .530 The minimum [frequency lor oy=27 15, The minimum frequency for $\phi_0 = 2n\pi$ is531Evolution of low-mass voung stellar objects (YSOs) can be traced by the infrared sul iin radio spectral enerev distributions (SEDs): molecular cloud cores through protostars are represented by Class 0 aud Class I SEDs. while classical T Tauri stars (CTTSs) and weak-lne T Tauri stars (ΝΤΤο) exhibit Class II aud Class III SEDs (Shu. Adams. Lizano 1987: André Moutiuerle 199D.,"Evolution of low-mass young stellar objects (YSOs) can be traced by the infrared – sub $mm$ radio spectral energy distributions (SEDs): molecular cloud cores through protostars are represented by Class 0 and Class I SEDs, while classical T Tauri stars (CTTSs) and weak-line T Tauri stars (WTTSs) exhibit Class II and Class III SEDs (Shu, Adams, Lizano 1987; André Montmerle 1994)."532 With the satelüte. T Tauri stars (TTSs = CTTSs aud WTTSs) have been found to cuit soft N-ravs. with the N-rav characteristics of moderate plasima temperature ( 1 keV). strong variability aud occasional rapid flares.," With the satellite, T Tauri stars (TTSs = CTTSs and WTTSs) have been found to emit soft X-rays, with the X-ray characteristics of moderate plasma temperature $\sim$ 1 keV), strong variability and occasional rapid flares."533 These are consistent with the scenario of eulianced solar-type activity. attributable to magnetic dynamo processes (Feieelsou DeCaipli 1981: Montiuerle et al.," These are consistent with the scenario of enhanced solar-type activity, attributable to magnetic dynamo processes (Feigelson DeCampli 1981; Montmerle et al."534 1983)., 1983).535 Protostars are often embedded deeply in star forming clouds. shrouded by dense circtuustellar eas aud dust. hence generally invisible in the optical. ucar infrared ancl even soft N-ray baucds.," Protostars are often embedded deeply in star forming clouds, shrouded by dense circumstellar gas and dust, hence generally invisible in the optical, near infrared and even soft X-ray bands."536 N-ravs from Class [sources (the late phase of protostars) have been reported iu deep audROSAT nuages of the p Ophiuchi cloud (INovama et al., X-rays from Class I sources (the late phase of protostars) have been reported in deep and images of the $\rho$ Ophiuchi cloud (Koyama et al.537 1991: Casanova oet al., 1994; Casanova et al.538 1995: Kamata et al., 1995; Kamata et al.539 1997: Grosso et al., 1997; Grosso et al.540 1997: Carkuer. Kozak. Feigelson 1998: Tsuboi 1999: Tsuboi et al.," 1997; Carkner, Kozak, Feigelson 1998; Tsuboi 1999; Tsuboi et al."541 2000: Cirosso et al., 2000; Grosso et al.542 2000: Crosso 2001)., 2000; Grosso 2001).543 ILowever possible contamination from nearby TTSs (or others) would uot be removed with the limited spatial resolution ofASCA., However possible contamination from nearby TTSs (or others) would not be removed with the limited spatial resolution of.544 Very low count rates from 20Shampered quantitative X-ray study ou these Class I sources., Very low count rates from hampered quantitative X-ray study on these Class I sources.545 Iu order to shed lieht ou the embedded YSOs. like ClassI sources in the p Ophiuchi cloud. we have observed this cloud. with theObservatory (CNO. Weisskopt. O/dell. vau Spevbroeck 1997). which has wide energy baud sensitivity coupled with uuprecedeuted spatial resolution.," In order to shed light on the embedded YSOs, like Class I sources in the $\rho$ Ophiuchi cloud, we have observed this cloud with the (CXO, Weisskopf, $'$ dell, van Speybroeck 1997), which has wide energy band sensitivity coupled with unprecedented spatial resolution."546 Iu this paper. we systematically study the X-ray features of the voung stellar objects in the p Ophiuchi cloud based on the Chandra results.," In this paper, we systematically study the X-ray features of the young stellar objects in the $\rho$ Ophiuchi cloud based on the $Chandra$ results."547 For this purpose. we first clarify our criteria of source classifications based on the near- to far-IR data.," For this purpose, we first clarify our criteria of source classifications based on the near- to far-IR data."548 Class I sources are selected from unear- to mid-IR baud or near- to far-IR baud spectra (André Moutinerle 1991: Casanova ct al., Class I sources are selected from near- to mid-IR band or near- to far-IR band spectra (André Montmerle 1994; Casanova et al.549 1995: Chen et al., 1995; Chen et al.550 1995: Chen et al., 1995; Chen et al.551 1997: Motto. Andzré.. Neri 1998: Luluunan Rieke 1999: Crosse et al.," 1997; Motte, André,, Neri 1998; Luhman Rieke 1999; Grosso et al."552 2000)., 2000).553 Since the classifications of these papers are not fully consistent with cach other. we regard a bona fide Class I source (Class I. here aud. after) if it is referred to as à Class [source by all the papers. aud the others to be Class E caucdidates (Class L.. hear aud after).," Since the classifications of these papers are not fully consistent with each other, we regard a bona fide Class I source (Class I, here and after) if it is referred to as a Class I source by all the papers, and the others to be Class I candidates (Class $_c$, hear and after)."554 Class IT aud III sources (here Class IT aud Class III) are simply selected by the spectral iudex between near- and nid- IR bauds (André Montinerle 1991: Casanova et al., Class II and III sources (here Class II and Class III) are simply selected by the spectral index between near- and mid- IR bands (André Montmerle 1994; Casanova et al.555 1995: Lubin Ricke 1999: Crosse et al., 1995; Luhman Rieke 1999; Grosso et al.556 2000)., 2000).557" The age or evolution sequence of Class 1, should be in between Class I aud Class IT.", The age or evolution sequence of Class $_c$ should be in between Class I and Class II.558" We also use he terminology ""Class ELLA by combining the Class I and Class L. for sinplicitv.", We also use the terminology “Class $_c$ ” by combining the Class I and Class $_c$ for simplicity.559" Likewise. ""Class IT|IIT ueans a conibined eroup of Class IT aud Class ΤΗ."," Likewise, “Class II+III” means a combined group of Class II and Class III."560 A brown dwarf (here aud after. BD) meas a star with mass falling well below the hvdrogeu-burniug iuit of 0.08 AL... and a brown dwarf candidate (bere and after. DD.) inclicates a star with mass close to 05 AL. (Wilkine. Greene. Alever 1999: Cushiug. Tokunaga. I&obavashi 2000).," A brown dwarf (here and after, BD) means a star with mass falling well below the hydrogen-burning limit of 0.08 $_{\odot}$, and a brown dwarf candidate (here and after, $_c$ ) indicates a star with mass close to 0.08 $_{\odot}$ (Wilking, Greene, Meyer 1999; Cushing, Tokunaga, Kobayashi 2000)."561 For comparison with the previous X-ray aud infrared results. we asstuue the distance to the p Oph region to be 165 pe (Dame et al.," For comparison with the previous X-ray and infrared results, we assume the distance to the $\rho$ Oph region to be 165 pc (Dame et al."562 1987) throughout this paper. although ITippareos data suggest a closer distance of 7.120 pe (suede Toe 1998).," 1987) throughout this paper, although $Hipparcos$ data suggest a closer distance of $\sim$ 120 pc (Knude Hog 1998)."563 The observation was made on 2000 April 1311 with the ACIS-I array consisting of four abutted Aaav CCDs. which covered a feld including the southeast part of the p ΟΡ cloud in the 0.210 keV baud.," The observation was made on 2000 April 13–14 with the ACIS-I array consisting of four abutted X-ray CCDs, which covered a field including the southeast part of the $\rho$ Oph cloud in the 0.2–10 keV band."564" The telescope optical axis position on the ACIS-T array is R.À. = 16527""18.p. Dec = 2321.9"" (epoch 2000)."," The telescope optical axis position on the ACIS-I array is R.A. = $^{\rm h}$ $^{\rm m}$ $^{\rm s}$, Dec = $-$ $^{\circ}$ $'$ $''$ (epoch 2000)."565 We use Level 1 processed events provided by the pipeline processing at the Chandra N-vav Center., We use Level 1 processed events provided by the pipeline processing at the Chandra X-ray Center.566 The total available exposure time is c 107 ks., The total available exposure time is $\approx$ $10^2$ ks.567 To minimize the effect of the degradation of charge trausfer inefiicicucy (CTI). we apply the improved data processing techuique developed by Townsley et al. (," To minimize the effect of the degradation of charge transfer inefficiency (CTI), we apply the improved data processing technique developed by Townsley et al. ("5682000).,2000).569 To reject backeround eveuts. we apply a grade filter to keep only 0. 2.3. L and 6.," To reject background events, we apply a grade filter to keep only 0, 2, 3, 4, and 6."570 Figure 1 shows the ACIS-I false-color image of the p Oph cloud., Figure 1 shows the ACIS-I false-color image of the $\rho$ Oph cloud.571 Red aud blue colors represent pliotous in the soft (0.7.2.0 keV) aud hard (2.07.0 keV) X- bands. respectively.," Red and blue colors represent photons in the soft (0.7–2.0 keV) and hard (2.0–7.0 keV) X-ray bands, respectively."572 For source finding. we run the," For source finding, we run the"573that their reconstructed. magnetic topologies have large regions of near-surface azimuthal field.,that their reconstructed magnetic topologies have large regions of near-surface azimuthal field.574 These are interpreted as the toroidal component of the large-scale dynamo field in the stars., These are interpreted as the toroidal component of the large-scale dynamo field in the stars.575 Phe presence of these regions near the stellar surface has led to the belief that. the dynamo operating in such stars is in [act distributed. throughout the stellar convective zone. rather than being restricted to the interface laver as in the solar case.," The presence of these regions near the stellar surface has led to the belief that the dynamo operating in such stars is in fact distributed throughout the stellar convective zone, rather than being restricted to the interface layer as in the solar case."576 Dilferential rotation is still thought to play a role in the generation of magnetic fields in these stars. however how the dyvnamos in these stars operates is not well understood.," Differential rotation is still thought to play a role in the generation of magnetic fields in these stars, however how the dynamos in these stars operates is not well understood."577 Most theoretical models are based on our knowledge of the solar dvnamo., Most theoretical models are based on our knowledge of the solar dynamo.578 The mocdels of Ixitehatinov&Riiciger(1999) wedict that the level of dilferential rotation on an carly-C star should be greater than that on a mid-Ix. star and that the level of surface differential rotation should decrease for stars with shorter rotational periods., The models of \citet{KitchatinovLL:1999} predict that the level of differential rotation on an early-G star should be greater than that on a mid-K star and that the level of surface differential rotation should decrease for stars with shorter rotational periods.579 While the models of Ixüker&Riiclieer(2005) also show that the level of differential rotation on a star should. be dependent upon its cllective temperature. (with hotter stars having higher levels of differential rotation) but. only weakly dependent on its rotation rate., While the models of \citet{KukerM:2005} also show that the level of differential rotation on a star should be dependent upon its effective temperature (with hotter stars having higher levels of differential rotation) but only weakly dependent on its rotation rate.580 These are predictions that we can now observationally test., These are predictions that we can now observationally test.581 There are several techniques we can use do observationally measure the level of dilferential. rotation on a star., There are several techniques we can use to observationally measure the level of differential rotation on a star.582 Direct. starspot tracking from multiple Doppler images of the surface features (Le.CollierCameron.Donati 2002).. eross-correlating two independent Doppler images (Le.Donati&CollierCameron1997).. incorporating a differential rotation law into the Doppler imaging process (i.c.Donatictal.2000:Petit.&CollierCameron 2002).. or by Fourier analvsis of stellar line profiles. (ic.Reiners&Schmitt2002. 2003).," Direct starspot tracking from multiple Doppler images of the surface features \citep*[i.e.][]{CameronAC:2002}, cross-correlating two independent Doppler images \citep[i.e.][]{DonatiJF:1997a}, incorporating a differential rotation law into the Doppler imaging process \citep*[i.e.][]{DonatiJF:2000,PetitP:2002}, or by Fourier analysis of stellar line profiles \citep[i.e.][]{ReinersA:2002, ReinersA:2003}."583. Combining cillerential rotation measurements of a number of voung solar-tvpe stars. found. using the Doppler imagine method. Barnesetal.(2005) found that the level of differential rotation increases with stellar temperature with early-G stars having higher levels of differential. rotation than lower-mass stars. in agreement with the findings of I|xitchatinov&Riiciger(1999)— and WKuker&ltüdiger (2005).," Combining differential rotation measurements of a number of young solar-type stars, found using the Doppler imaging method, \citet{BarnesJR:2005} found that the level of differential rotation increases with stellar temperature with early-G stars having higher levels of differential rotation than lower-mass stars, in agreement with the findings of \citet{KitchatinovLL:1999} and \citet{KukerM:2005}."584. The results also showed. only a weak (if anv) correlation with stellar rotation rate. again in agreement with the findings of Ixüker&Itücdiger(2005).," The results also showed only a weak (if any) correlation with stellar rotation rate, again in agreement with the findings of \citet{KukerM:2005}."585. Llowever. more recent measurements of the voung earlv-G star HD. 171488 (Miusdenetal.2006:Jellers&Donati2008:Jelfoersοἱ2010) have shown this star to have an even higher level ofcülferential rotation than that indicated by Barnesοἱal.(2t5) with dillerential rotation measurements up to 10 times the current solar value.," However, more recent measurements of the young early-G star HD 171488 \citep{MarsdenSC:2006, JeffersSV:2008, JeffersSV:2010} have shown this star to have an even higher level of differential rotation than that indicated by \citet{BarnesJR:2005} with differential rotation measurements up to 10 times the current solar value."586 This supports the findings of Reiners(2006) which show high levels of dilferential rotation on à number of E stars., This supports the findings of \citet{ReinersA:2006} which show high levels of differential rotation on a number of F stars.587 Work bv Donati.CollierCameron&Petit(2003b) has shown that for some earlv-Ix. stars the level of dillerential rotation measured from surface brightness [features is consistently lower than that measured [rom the magnetic features (using the same dataset)., Work by \citet*{DonatiJF:2003b} has shown that for some early-K stars the level of differential rotation measured from surface brightness features is consistently lower than that measured from the magnetic features (using the same dataset).588 This has led them to surmise that this is a result of the magnetic ancl brightness features being anchored at cdilferent depths within the stellar convective zone and that the convective zone has a raciallv varving dillerential rotation structure. unlike the Sun.," This has led them to surmise that this is a result of the magnetic and brightness features being anchored at different depths within the stellar convective zone and that the convective zone has a radially varying differential rotation structure, unlike the Sun."589 The work also showed that these earlbv-I stars evidence temporal variation in their levels of cülferential rotation that the authors attribute to a feed-back mechanism in the stellar dynamo periodically converting magnetic energy into kinetic energy ancl vice-versa., The work also showed that these early-K stars evidence temporal variation in their levels of differential rotation that the authors attribute to a feed-back mechanism in the stellar dynamo periodically converting magnetic energy into kinetic energy and vice-versa.590 As mentioned there is currently only one carly-G star or which there are dillerential rotation measures [roni xh brightness and magnetic features obtained at multiple epochs. LID 171488 (Marsdenetal.2006:Jelfers&Do-nati2008:Jeffersetal. 2010).," As mentioned there is currently only one early-G star for which there are differential rotation measures from both brightness and magnetic features obtained at multiple epochs, HD 171488 \citep{MarsdenSC:2006, JeffersSV:2008, JeffersSV:2010}."591.. This work has shown very ittle difference in the level of dillerential rotation measured rom the magnetic and brightness features. although the nmieasurement errors are larger than that found on the carly-Ix stars (Donatietal.2003b).," This work has shown very little difference in the level of differential rotation measured from the magnetic and brightness features, although the measurement errors are larger than that found on the early-K stars \citep{DonatiJF:2003b}."592. Xcdditionallv. there appears o be little evidence of temporal evolution in the level of differential rotation on HD 1712488.," Additionally, there appears to be little evidence of temporal evolution in the level of differential rotation on HD 171488."593 This has been speculated o be caused by the thinner convective zone of HD. 171488 compared to that of the earlv-Ix stars previously. stucliect., This has been speculated to be caused by the thinner convective zone of HD 171488 compared to that of the early-K stars previously studied.594 In order to expand the number of voung earlv-Ci stars studied: with spectropolarimetry this paper along with the irst paper in the series Marsdenetal.(2010.Paper1) and Waiteet.al.(2010)... presents dillerential. rotation measurements. and magnetic maps. of two voung carly-C oe-main sequence (PAIS) stars.," In order to expand the number of young early-G stars studied with spectropolarimetry this paper along with the first paper in the series \citet[][Paper I]{MarsdenSC:2010} and \citet{WaiteIA:2010}, presents differential rotation measurements, and magnetic maps, of two young early-G pre-main sequence (PMS) stars."595 Paper E deals with the reconstruction of the brightness ancl magnetic topologies of the star LID 141943., Paper I deals with the reconstruction of the brightness and magnetic topologies of the star HD 141943.596 ‘This paper Paper LL) presents he coronal magnetic Ποια reconstructions. lla activity and. dillerential rotation measurements of LD 141943.," This paper (Paper II) presents the coronal magnetic field reconstructions, $\alpha$ activity and differential rotation measurements of HD 141943."597 The hired paper in the series (Waitectal.2010) eals with observations of LID 106506., The third paper in the series \citep{WaiteIA:2010} deals with observations of HD 106506.598 As detailed in Paper Ll. HD. 141948 is a voung. active ancl rapidly-rotating PALS star.," As detailed in Paper I, HD 141943 is a young, active and rapidly-rotating PMS star."599 Ehe stellar parameters have been determined in Paper 1 through the Doppler imaging process and are reproduced here in Table 1L.., The stellar parameters have been determined in Paper I through the Doppler imaging process and are reproduced here in Table \ref{Tab_par}.600 The reconstructed. brightness images show that it has a weak polar spot and a significant amount. of low-latituce spot features at. all epochs., The reconstructed brightness images show that it has a weak polar spot and a significant amount of low-latitude spot features at all epochs.601 Phe magnetic reconstructions show that its has a predominately non-axisvmametric radial field while its azimuthal field is predominately axisvnunietric with ring of azimuthal field seen at the pole. similar to that of other active stars.," The magnetic reconstructions show that its has a predominately non-axisymmetric radial field while its azimuthal field is predominately axisymmetric with ring of azimuthal field seen at the pole, similar to that of other active stars."602 11D141043 was observed at 3 epochs on the 8.9-n1 Anglo- telescope (AAT). in March/April 2007. April 2009 ancl March/April 2010.," HD141943 was observed at 3 epochs on the 3.9-m Anglo-Australian telescope (AAT), in March/April 2007, April 2009 and March/April 2010."603 A fourth observation set was, A fourth observation set was604of the new pair of images identified in this letter (5.1 / 5.2) is in excellent agreement (within 0.5”) with the predictions from this model at z=6.,of the new pair of images identified in this letter (5.1 / 5.2) is in excellent agreement (within ) with the predictions from this model at $z=6$.605 We therefore use this new information as an additional constraint (system 5) to improve the mass model., We therefore use this new information as an additional constraint (system 5) to improve the mass model.606" From the family of best models resulting from the optimisation, we derive a magnification factor µι=11.4t1.9 and we=7.3+1.2 for images 5.1 and 5.2, respectively."," From the family of best models resulting from the optimisation, we derive a magnification factor $\mu_1=11.4\pm1.9$ and $\mu_2= 7.3\pm1.2$ for images 5.1 and 5.2, respectively."607 The ratio µι/µο is in perfect agreement with the F110W photometry., The ratio $\mu_1/\mu_2$ is in perfect agreement with the F110W photometry.608" A third image (5.3) is predicted but is demagnified and located at the very centre of the BCG, making its detection extremely challenging, if not impossible."," A third image (5.3) is predicted but is demagnified and located at the very centre of the BCG, making its detection extremely challenging, if not impossible."609" Taking into account these magnification factors, we derive intrinsic properties for the original background source 5."," Taking into account these magnification factors, we derive intrinsic properties for the original background source 5."610" Its unlensed F110W (J band) magnitude (27.4 AB) corresponds to an ~0.4L* galaxy at z=6, which makes it more representative of the galaxy population at this redshift than brighter objects, such as the i-dropout identified in Abell 1703 (?) which corresponds to a ~3.6L* galaxy."," Its unlensed F110W (J band) magnitude (27.4 AB) corresponds to an $\sim$ $L^*$ galaxy at $z=6$, which makes it more representative of the galaxy population at this redshift than brighter objects, such as the i-dropout identified in Abell 1703 \citep{Richard09} which corresponds to a $\sim$ $L^*$ galaxy."611" A particularly interesting feature of source 5 is the large magnitude break (1.50+0.10 mag) between the H band (F160W) and the two first channels of IRAC, passbands that straddle the location of the rest-frame Balmer break."," A particularly interesting feature of source 5 is the large magnitude break $1.50\pm0.10$ mag) between the H band (F160W) and the two first channels of IRAC, passbands that straddle the location of the rest-frame Balmer break."612" In order to interpret the implications of this break for the properties of the underlying stellar populations, we performed a Spectral Energy Distribution (SED) fit of image 5.1 (highest S/N image) with a library of spectral templates following the precepts discussed in detail by ?.."," In order to interpret the implications of this break for the properties of the underlying stellar populations, we performed a Spectral Energy Distribution (SED) fit of image 5.1 (highest S/N image) with a library of spectral templates following the precepts discussed in detail by \citet{Stark09}."613" The main parameters we aim to derive are the total stellar mass (M*), the current star-formation rate (SFR) and the age of the stellar population (T)."," The main parameters we aim to derive are the total stellar mass $^*$ ), the current star-formation rate (SFR) and the age of the stellar population $T$ )."614" To this end, we removed the contribution of the measured Lyman-o line in the F814W and F850LP photometry, and then fit Charlot Bruzual(2007, S. Charlot, private communication) stellar- synthesis models to the observed SEDs."," To this end, we removed the contribution of the measured $\alpha$ line in the F814W and F850LP photometry, and then fit Charlot Bruzual(2007, S. Charlot, private communication) stellar-population synthesis models to the observed SEDs."615" We considered exponentially decaying star formation histories of the form SFR(t)~exp(—t/7) with e-folding times of T—10, 70, 100, 300, and 500 Myr, in addition to models with continuous star formation (CSF)."," We considered exponentially decaying star formation histories of the form $\rm{SFR(t)}\simeq exp(-t/\tau)$ with e-folding times of $\tau=10$, 70, 100, 300, and 500 Myr, in addition to models with continuous star formation (CSF)."616 T' is restricted to the range of 10 to 1000 Myr (the age of the universe at z~ 6)., $T$ is restricted to the range of 10 to 1000 Myr (the age of the universe at $z\sim6$ ).617 We use a Salpeter initial mass function (IMF) (?) and the dust extinction law of ? with 0.0«E(B—V)<1.0., We use a Salpeter initial mass function (IMF) \citep{Salpeter} and the dust extinction law of \citet{Calzetti} with $0.0<E(B-V)<1.0$.618" Finally, we allow the metallicity to vary between solar (Zo) and 0.2 Zo, but note that metallicity is degenerate with reddening and has a second-order effect on the SED."," Finally, we allow the metallicity to vary between solar $_\odot$ ) and 0.2 $_\odot$ , but note that metallicity is degenerate with reddening and has a second-order effect on the SED."619 Absorption by the intergalactic medium is accounted for following ?.., Absorption by the intergalactic medium is accounted for following \citet{Meiksin}.620 We next consider the possibility that the strong break between the H band and the IRAC channels is due to contamination by nebular emission lines., .We next consider the possibility that the strong break between the H band and the IRAC channels is due to contamination by nebular emission lines.621" Strong Ho, H6 and [OIII| lines would all fall into the IRAC bands and could boost the continuum, mimicking a stronger Balmer break (?).."," Strong $\alpha$, $\beta$ and $[OIII]$ lines would all fall into the IRAC bands and could boost the continuum, mimicking a stronger Balmer break \citep{Schaerer10}."622" We correct for this effect as follows: the SFR estimated from SED fitting is consistent with the observed Lyman-a flux, when assuming an aperture correction of a factor of 2 and (conservatively) 1/3 case B for Lyman-a."," We correct for this effect as follows: the SFR estimated from SED fitting is consistent with the observed $\alpha$ flux, when assuming an aperture correction of a factor of 2 and (conservatively) 1/3 case B for $\alpha$."623" We predict observed fluxes faa=1.46 10!"", fgg=5.1 1075 and =6.7 107? (cgs units, both [OIII] lines average fromfior COSMOS galaxies, Zoubian et al."," We predict observed fluxes $f_{H\alpha}= 1.46$ $10^{-17}$, $f_{H\beta}=5.1$ $10^{-18}$ and $f_{\rm[OIII]}=6.7$ $10^{-18}$ (cgs units, both [OIII] lines average from COSMOS galaxies, Zoubian et al."624 2011 in prep.), 2011 in prep.).625" We therefore lower the fluxes by 0.06 and 0.04 magnitudes at 3.6 wm and 4.5jum, respectively, thereby eliminating contamination from unresolved line emission."," We therefore lower the fluxes by 0.06 and 0.04 magnitudes at 3.6 $\mu$ m and $\mu$ m, respectively, thereby eliminating contamination from unresolved line emission."626 This correction is negligible and we observe no significant change in the best-fit model using this updated photometry., This correction is negligible and we observe no significant change in the best-fit model using this updated photometry.627 The observed SED and the best-fitting template are presented in Fig. 4.., The observed SED and the best-fitting template are presented in Fig. \ref{sed}. .628 The best fit is achieved for a 7—500 Myr star-formation history., The best fit is achieved for a $\tau=500$ Myr star-formation history.629" The full set of templates gives an age range of T=640-940 Myrs (within a Ay?=3 from the best model), depending on the exact star-formation history, corresponding to a redshift of formation of z;=18+4. More punctuated star-formation histories (r—10100 Myrs) could produce similar Balmer breaks with ages of onlyTT—130300 Myrs, but they produce much poorer fits to thedata (Ax?~50, 4)). (?) ?)) ?.. (?).."," The full set of templates gives an age range of $T$ =640-940 Myrs (within a $\Delta\chi^2=3$ from the best model), depending on the exact star-formation history, corresponding to a redshift of formation of $z_f=18\pm4$ More punctuated star-formation histories $\tau=10-100$ Myrs) could produce similar Balmer breaks with ages of only$T=130-300$ Myrs, but they produce much poorer fits to thedata $\Delta\chi^2\sim50$ \ref{sed}) \citep{Poggianti} \citealt{Brinchmann2004}) \citet{Egami}. \citep{Bouwens_slope}."630"favorable for =2-land, y is more favorable for lJI",favorable for $=2-4$and $\tau_{\it eff}^{-}$ is more favorable for $>4$.631ITowever. wedonotiusistthatthistrendisrcaleonsideringlli Ισ Lal y vofgala," However, we do not insist that this trend is real considering the fact that we do not have much knowledge about the UV flux of galaxies."632a noc aat22x Lby carefully treating the intergalactic Lii absorption.," Nonetheless, we can obtain more accurate at $2<z<4$ by carefully treating the intergalactic Lyman absorption."633 Iu sunuuuw. we find that the internal absorption is effective iu reducing the catastrophic errors. while the iutergalactie absorption is effective in reducing the systelmatic error af 2—2.LK.," In summary, we find that the internal absorption is effective in reducing the catastrophic errors, while the intergalactic absorption is effective in reducing the systematic error at $z=2-4$."634 We have used 187 template SEDs covering a wide rauge of spectral types., We have used 187 template SEDs covering a wide range of spectral types.635 Generally speaking. accuratephoto. even for high-: galaxies. can be obtained using either observed spectra (CWW: Isinmey et al.," Generally speaking, accurate, even for $z$ galaxies, can be obtained using either observed spectra (CWW; Kinney et al."636 1993) or simulated spectra (NAOT: BC96). though absorption effects Guternal and intergalactic absorption) must be carefully treated.," 1993) or simulated spectra (KA97; BC96), though absorption effects (internal and intergalactic absorption) must be carefully treated."637 If the purpose of the analysis is just to obtain ((e.g.. Ferndundez-Soto et al.," If the purpose of the analysis is just to obtain (e.g., Fernánndez-Soto et al."638 1999). a few template SEDs would be suf&cieut.," 1999), a few template SEDs would be sufficient."639 However. if we want to determine accurately the SED shape of a target galaxy in order to derive useful information ou the SED. wach more template SEDs are necessary as we will see below.," However, if we want to determine accurately the SED shape of a target galaxy in order to derive useful information on the SED, much more template SEDs are necessary as we will see below."640" Figure 22 shows oof the TDF ealaxics using four SEDs from CWW as cluplate SEDs, which is almost the same procedure as Feruáunudez-Soto et al."," Figure \ref{fig:photoz_CWW4} shows of the HDF galaxies using four SEDs from CWW as template SEDs, which is almost the same procedure as Fernánndez-Soto et al."641 adopted (Note that we reproduce he same trend in Fieure 22 as secu in Figure 19((0))., adopted (Note that we reproduce the same trend in Figure \ref{fig:photoz_CWW4} as seen in Figure \ref{fig:comp}( (c)).642 The accuracy Of ybased on the four SEDs is worse than that of our best estimate (Figure 19((0)). but the differeuce is within a actor of about 2.," The accuracy of based on the four SEDs is worse than that of our best estimate (Figure \ref{fig:comp}( (c)), but the difference is within a factor of about 2."643 IIowever. there is a large difference iu lio minium reducedν," However, there is a large difference in the minimum reduced."644"τι, Our minium reduced vvalues are plotted in Figure 6..", Our minimum reduced values are plotted in Figure \ref{fig:chi2}.645 Figure 223 is thle same lot as Figure 6. but for the four-SED case.," Figure \ref{fig:chi2_CWW4} is the same plot as Figure \ref{fig:chi2}646 but for the four-SED case."647 The minimal reduced vvalues obtained with the four-SED templates are much arecr than those with 3.366 (187«GE(DVW)\ 37) emplates.," The minimum reduced values obtained with the four-SED templates are much larger than those with 3,366 $187\times 6 E(B-V) \times 3\tau $ ) templates."648 The distribution of reduced vvalues can be a measure to determine how mauy template SEDs are necessary aud sufficicut for a given purpose., The distribution of reduced values can be a measure to determine how many template SEDs are necessary and sufficient for a given purpose.649 We obtain photometric redshifts of 925 ealaxics out of the 916 ealaxies in the good S/N region (See 77))., We obtain photometric redshifts of 925 galaxies out of the 946 galaxies in the good S/N region (See \ref{sec:data}) ).650 We call these 925 ealaxies the plotometric sauple., We call these 925 galaxies the photometric sample.651 The redshift distribution .N(:) of the plotometric sample is shown in Fieure 21. , The redshift distribution $N(z)$ of the photometric sample is shown in Figure \ref{fig:compare_nz}. .652For comparison. the results bv Fornaundez-Soto et al.," For comparison, the results by Fernánndez-Soto et al."653 are superposed., are superposed.654 Sawicki et al. (, Sawicki et al. (6551997) argued that Ας) could be an indicator of the error iu photometric redshift technique.,1997) argued that $N(z)$ could be an indicator of the error in photometric redshift technique.656 They claimed thatvestry among spectral brealss 000A... Bahuer.2800A.. and 2635À)) easilv leads to au uurealistic ANC) which las typically two pronunent peaks at :20.1 and zz 2.," They claimed that among spectral breaks , Balmer, and ) easily leads to an unrealistic $N(z)$, which has typically two prominent peaks at $z\simeq 0-1$ and $z\simeq 2$ ."657 The peaks are likely caused by a template set which does not include the internal or the iuterealactic absorption., The peaks are likely caused by a template set which does not include the internal or the intergalactic absorption.658 Our ;V(:) does not have such two prominent peaks., Our $N(z)$ does not have such two prominent peaks.659 Iustead. our IN(:) shows a single peak at £220.5.1 and a immoderate decrease at 21 at all maguitude bius.," Instead, our $N(z)$ shows a single peak at $z\simeq 0.5-1$ and a moderate decrease at $z>1$ at all magnitude bins."660 Foruáundez-Soto et al, Fernánndez-Soto et al.661 found that ealaxics at :=ο ον wnnt (Figure 19((00)), found that galaxies at $z=2-4$ have systematically lower photometric redshfits (Figure \ref{fig:comp}( (c)).662 We find in this paper that this is attributed to a systematic error which is sensitive to the intergalactic absorption., We find in this paper that this is attributed to a systematic error which is sensitive to the intergalactic absorption.663 This ‘small aliasine’ due to the systematic error would make au uulialistie peak at; 2 in τὸν especially for faint galaxies.," This `small aliasing' due to the systematic error would make an unlialistic peak at $z\simeq 2$ in $N(z)$, especially for faint galaxies."664 For 26<I«28. the AN) of Fernánudez-Soto et al.," For $26<I<28$, the $N(z)$ of Fernánndez-Soto et al."665 has a prominent peak at22., has a prominent peak at$z\simeq 2$.666 This is probably due to the small aliasing of galaxies at z22. because the wmmber of Fernauucdez-Soto et al.," This is probably due to the small aliasing of galaxies at $z\ge 2$, because the number of Fernánndez-Soto et al."667's ealaxies with :>2 is smaller than that of our galaxies.,'s galaxies with $z>2$ is smaller than that of our galaxies.668 Next. we examine £(2BVW) of galaxies.," Next, we examine $E(B-V)$ of galaxies."669 Figure 25 shows the histograms of E(BV) of the best-fit SEDs in the five redshift rauges., Figure \ref{fig:dust_z} shows the histograms of $E(B-V)$ of the best-fit SEDs in the five redshift ranges.670 The arrow in cach redshift ranec dicates the mean value., The arrow in each redshift range indicates the mean value.671 We find that the E(BY) value of galaxies remains almost constant (E(BV)=0.13.0.18) for το6. though it has a simall poak at 2=ὉL bin.," We find that the $E(B-V)$ value of galaxies remains almost constant $E(B-V)=0.13-0.18$ ) for $z=0-6$, though it has a small peak at $z=3-4$ bin."672 These values are roughly cousisteut with those of Lyiman-break ealaxies at 2=2.53.5 derived bv Sawicki Yee (1998: E(BΤσ0.3). because our photometric sample probably includes uot ouly star-forming Lyuiur-break galaxies but also passively evolving elliptical galaxies.," These values are roughly consistent with those of Lyman-break galaxies at $z=2.5-3.5$ derived by Sawicki Yee (1998; $E(B-V)\simeq 0.3$ ), because our photometric sample probably includes not only star-forming Lyman-break galaxies but also passively evolving elliptical galaxies."673 Calzetti Heckman (1999) iteratively calculated the cosmic star formation rate (SFR) aud the mean £(B V)of galaxies using the Calactic conversion factor between the gas column deusity {1} aud E(BVine. BOBV)=NUTD/5.9«102 (Bohlin. Savage. Drake 1978).," Calzetti Heckman (1999) iteratively calculated the cosmic star formation rate (SFR) and the mean $E(B-V)$ of galaxies using the Galactic conversion factor between the gas column density $N(H)$ and $E(B-V)$ , i.e., $E(B-V)=N(H)/5.9\times 10^{21}$ (Bohlin, Savage, Drake 1978)."674 The vertical lines plotted im Figure 25 are predictions of E(BV) based on their two typical models., The vertical lines plotted in Figure \ref{fig:dust_z} are predictions of $E(B-V)$ based on their two typical models.675 The model A corresponds to the cosimic SER which remains almost constaut from :=1.5 to 5., The model `A' corresponds to the cosmic SFR which remains almost constant from $z=1.5$ to $5$.676 The model D corresponds to the cosmic SER which reaches its peak at >=12 anc monotonically decreases at +=25., The model `B' corresponds to the cosmic SFR which reaches its peak at $z=1-2$ and monotonically decreases at $z=2-5$.677 Qur results ave found to be more consistent with the model AD., Our results are found to be more consistent with the model `A'.678 We improve photometric redshifts of LOLS galaxies iu the ΠΟΕ., We improve photometric redshifts of 1048 galaxies in the HDF.679 A standard nuuimimuizing method is adopted for obtainiusgp, A standard minimizing method is adopted for obtaining.680hotos We use 187 template SEDs represcuting a wide variety of morphologv aud age of observed galaxies based ou a population svuthesis model by Kodama Arimoto (1997) with two new recipes., We use 187 template SEDs representing a wide variety of morphology and age of observed galaxies based on a population synthesis model by Kodama Arimoto (1997) with two new recipes.681 First. the amount of the imternal absorption is changed as a free parameter in the ranee of E(BVV) =0.0to 0.5 with an interval of 0.1.," First, the amount of the internal absorption is changed as a free parameter in the range of $E(B-V)=0.0$ to $0.5$ with an interval of $0.1$."682 Secoud. the absorption due to interealactic IIT clouds is also chauged by a factor of 0.5. 1.0. aud 1.5 around the opacity given by AMadau (1995).," Second, the absorption due to intergalactic HI clouds is also changed by a factor of 0.5, 1.0, and 1.5 around the opacity given by Madau (1995)."683 The total ummber of template SEDs is thus 187«6«3=3.366. except for theredshift ον].," The total number of template SEDs is thus $187\times 6\times 3=3,366$, except for theredshift grid."684 The introduction ofthe internal absorption is ford to be effective in reducing the catastrophic errors i. wwhile the absorption due to the intergalactic HIT clouds helps to improve the svstematic error nn aat 2= l1., The introduction of the internal absorption is found to be effective in reducing the catastrophic errors in while the absorption due to the intergalactic HI clouds helps to improve the systematic error in at $z=2-4$ .685" The dispersion o. of our photometric redshiftswith respectto spectroscopicredshifts is σ.= and 0.21 for 2< and :7 2. respectively, which are smaller than the correspoucding values (o.=0.09 and 0.15) by Ferndundez-Soto et al. ("," The dispersion $\sigma_z$ of our photometric redshiftswith respectto spectroscopicredshifts is $\sigma_z=0.08$ and $0.24$ for $z<2$ and $z>2$ , respectively, which are smaller than the corresponding values $\sigma_z=0.09$ and $0.45$ ) by Fernánndez-Soto et al. ("6861999).,1999).687 Siguificaut inrrovenients are obtained. especially at hieh-:.," Significant improvements are obtained, especially at $z$ ."688 A conrparisoun with previous work and the properties of the best-fit SEDs for the spectroscopic sample verify the validity of our photometric redshifts., A comparison with previous work and the properties of the best-fit SEDs for the spectroscopic sample verify the validity of our photometric redshifts.689"the uncertainty in our methods is dominated mostly by the low number of reference sources above 4.85 GHz, we expect a better agreement as the pool of reference sources increases.","the uncertainty in our methods is dominated mostly by the low number of reference sources above $4.85$ GHz, we expect a better agreement as the pool of reference sources increases."690" In the range 0.5—1.0 Jy, simulations show evident artifacts due to the transition between different regimes of flux extrapolation, i.e. using spectral index from the WMAPS catalog for S23>1 Jy."," In the range $0.5 -1.0$ Jy, simulations show evident artifacts due to the transition between different regimes of flux extrapolation, i.e. using spectral index from the WMAP5 catalog for $S_{23} > 1$ Jy."691" They are due to the low number of sources in that flux range in the base catalogs; as above, we expect these artifact to significantly decrease as more data over the relevant frequency and flux range will be included in the reference catalogs."," They are due to the low number of sources in that flux range in the base catalogs; as above, we expect these artifact to significantly decrease as more data over the relevant frequency and flux range will be included in the reference catalogs."692" While our extrapolations correctly recover WMAPS results at lower frequencies, counts at 94 GHz are overestimated by ~50%."," While our extrapolations correctly recover WMAP5 results at lower frequencies, counts at 94 GHz are overestimated by $\sim 50\%$."693 In figure 6 we plot the integrated counts N(>S=1Jy) from our simulations and from WMAPS data.," In figure \ref{fig:int_counts} we plot the integrated counts $N(>S=1{\rm694 Jy})$ from our simulations and from WMAP5 data."695" For reference, we also show an analytic estimate of the counts assuming dN/dS«ο?? and Scv99? (Wrightetal.2009), normalized to fit WMAP5 counts in K, Ka and Q bands."," For reference, we also show an analytic estimate of the counts assuming $dN/dS \propto696S^{-2.5}$ and $S\propto \nu^{-0.09}$ \citep{2009ApJS..180..283W}, normalized to fit WMAP5 counts in K, Ka and Q bands."697" WMAPS data show steepening of the counts in W band compared to lower frequencies,a while our counts seem to be in better agreement with a linear log(S)—log(v) relation."," WMAP5 data show a steepening of the counts in W band compared to lower frequencies, while our counts seem to be in better agreement with a linear $\log(S)698- \log(\nu)$ relation."699" As discussed in section ??,, our extrapolations of sources with S23>1 Jy is mainly based on the family of spectral indexes provided by the WMAPS team, which can be approximated by a Gaussian distribution with mean a=—0.09 and dispersion σα=0.17."," As discussed in section \ref{sec:methods}, our extrapolations of sources with $S_{23} > 1$ Jy is mainly based on the family of spectral indexes provided by the WMAP5 team, which can be approximated by a Gaussian distribution with mean $\alpha = -0.09$ and dispersion $\sigma_\alpha = 0.17$."700" Therefore, as expected source counts in the simulations are in agreement with the analytic estimate plotted in figure 6;; the slightly steeper behaviour seen in simulations is compatible with the actual distribution of WMAPS spectral indexes having a tail toward steeper a."," Therefore, as expected source counts in the simulations are in agreement with the analytic estimate plotted in figure \ref{fig:int_counts}; the slightly steeper behaviour seen in simulations is compatible with the actual distribution of WMAP5 spectral indexes having a tail toward steeper $\alpha$."701 This suggests that there is some tension between the WMAPS5 W-band counts and their spectral index distribution., This suggests that there is some tension between the WMAP5 W-band counts and their spectral index distribution.702 A possible explanation could be a progressive steepening of the spectral index of bright sources with increasing frequency., A possible explanation could be a progressive steepening of the spectral index of bright sources with increasing frequency.703" González-Nuevoetal.(2008) noticed that the spectral index of WMAP sources in the 41—64 GHz range is steeper than their spectral index in the 5— 23GHz range, and radio source measurements in the ~100—250 GHz range by the QUaD telescope (Friedmanetal.2009) and the South Pole Telescope (SPT,Vieiraetal.2009) provide additional evidence in this direction (seedeZottietal.2009,forarecentreview).."," \cite{2008MNRAS.384..711G} noticed that the spectral index of WMAP sources in the $41-64$ GHz range is steeper than their spectral index in the $5-23$ GHz range, and radio source measurements in the $\sim 100 -250$ GHz range by the QUaD telescope \citep{2009ApJ...700L.187F} and the South Pole Telescope \citep[SPT,][]{2009arXiv0912.2338V} provide additional evidence in this direction \citep[see][for a recent review]{2009A&ARv.tmp...15D}."704" If this is actually the case, describing the source behaviour in the 23—94 GHz range with a single power law may not be entirely accurate."," If this is actually the case, describing the source behaviour in the $23 -94$ GHz range with a single power law may not be entirely accurate."705" Alternatively, there may be some unaccounted for systematics affecting WMAP detection efficiency in W band."," Alternatively, there may be some unaccounted for systematics affecting WMAP detection efficiency in W band."706" The WMAP team required that for a source to be included in the catalog it needed to be detected at more than 5c, in at least one band."," The WMAP team required that for a source to be included in the catalog it needed to be detected at more than $5\sigma$, in at least one band."707" Its flux in the other bands would be included if: a) it was measured at more than 2c, b) the fitted source width is within a factor of 2 of the actual beam (Wrightetal.2009).. Sawa"," Its flux in the other bands would be included if: a) it was measured at more than $2\sigma$, b) the fitted source width is within a factor of 2 of the actual beam \citep{2009ApJS..180..283W}."708ngwit&Shanks(2009) pointed out that the W band beam response to discrete sources appears to depend in a non-linear fashion on source flux., \cite{2009arXiv0912.0524S} pointed out that the W band beam response to discrete sources appears to depend in a non-linear fashion on source flux.709" In particular, they find that the effective beam profile is significantly wider for point sources than for the Jupiter observations used as a basis for the WMAP team's beam analysis, which can create a significant bias in determining if a source fits requirement (b)."," In particular, they find that the effective beam profile is significantly wider for point sources than for the Jupiter observations used as a basis for the WMAP team's beam analysis, which can create a significant bias in determining if a source fits requirement (b)."710" Since number counts from simulations do not include systematics due to WMAP actual detection procedure, this effect, if confirmed, could be an alternative explanation for the discrepancies between simulated and actual number counts."," Since number counts from simulations do not include systematics due to WMAP actual detection procedure, this effect, if confirmed, could be an alternative explanation for the discrepancies between simulated and actual number counts."711" Further data on source behaviour at high frequency, e.g. from Planck, will help solve this issue."," Further data on source behaviour at high frequency, e.g. from Planck, will help solve this issue."712line-of-sight of +100 kpe is considered. the resulting volume. which is covered by the survey. is only about two percent of the volume of a sphere of radius 100 kpe. wherein the ECs would be distributed.,"line-of-sight of $\pm$ 100 kpc is considered, the resulting volume, which is covered by the survey, is only about two percent of the volume of a sphere of radius 100 kpc, wherein the ECs would be distributed."713 In addition. Larsen&Brodie(2000) and Larsenetal.(2002) have discovered a population of ECs co-rotating with the disk of the lenticular galaxy NGC1023.," In addition, \cite{larbro00} and \cite{larsen02} have discovered a population of ECs co-rotating with the disk of the lenticular galaxy NGC1023."714 These so-called faint fuzzies have similar structural parameters as halo ECs and are therefore not easily distinguishable from halo ECs projected onto the disk on the basis of imaging data alone., These so-called faint fuzzies have similar structural parameters as halo ECs and are therefore not easily distinguishable from halo ECs projected onto the disk on the basis of imaging data alone.715 A fair fraction of ECs found in extragalactic surveys might therefore be associated with the disks and not the halos of these galaxies.Burkertetal.(2005) analyzed the kinematics of the faint fuzzies and concluded that they form a ring-like structure within the galactic disk of 11023 and that this ring was probably formed during a galaxy-galaxy interaction comparable to the Cartwheel galaxy., A fair fraction of ECs found in extragalactic surveys might therefore be associated with the disks and not the halos of these \cite{burkert} analyzed the kinematics of the faint fuzzies and concluded that they form a ring-like structure within the galactic disk of 1023 and that this ring was probably formed during a galaxy-galaxy interaction comparable to the Cartwheel galaxy.716 A detailed discussion of faint fuzzies in the context of merged CCs is given in Brüns (2009).. who demonstrated that the observed structural parameters of the faint fuzzies are in excellent agreement with the merged CC scenario.," A detailed discussion of faint fuzzies in the context of merged CCs is given in \cite{bruens09}, who demonstrated that the observed structural parameters of the faint fuzzies are in excellent agreement with the merged CC scenario."717 Another reason for incompleteness is the difficulty of distinguishing ECs from background galaxies., Another reason for incompleteness is the difficulty of distinguishing ECs from background galaxies.718 The GC surveys covering 100 galaxies of the Virgo Cluster (Jordánetal.2005) and 43 galaxies of the Fornax Cluster (Mastersetal.2010) applied a size limit of raj« 10 pe to reduce the contamination of background galaxies., The GC surveys covering 100 galaxies of the Virgo Cluster \citep{jordan05} and 43 galaxies of the Fornax Cluster \citep{masters10} applied a size limit of $r_{\rm eff} <$ 10 pc to reduce the contamination of background galaxies.719 Thereby. they excluded also all ECs from their GC catalogs.," Thereby, they excluded also all ECs from their GC catalogs."720 Pengetal.(2006) used the same Virgo Cluster survey data as Jordanetal.(2005) to search for diffuse star clusters and found e.g. in the galaxy about 30 ECs. where Jordanetal.(2005) found 211 compact GCs.," \cite{peng} used the same Virgo Cluster survey data as \cite{jordan05} to search for diffuse star clusters and found e.g. in the galaxy about 30 ECs, where \cite{jordan05} found 211 compact GCs."721 Pengetal.(2006) have demonstrated that hundreds of ECs await detection in galaxy clusters., \cite{peng} have demonstrated that hundreds of ECs await detection in galaxy clusters.722 Without follow-on spectroscopy to determine radial velocities. it cannot be decided whether these ECs are associated with the main body or the halo of the galaxies.," Without follow-on spectroscopy to determine radial velocities, it cannot be decided whether these ECs are associated with the main body or the halo of the galaxies."723 While lower-mass ECs up to Mpc=10° Me are ofter overlooked in surveys. much effort has been made to detect anc to analyze stellar objects of M=107 Me since the discovery of UCDs in the Fornax Cluster by Hilkeretal.(1999). anc Drinkwateretal.(2000).," While lower-mass ECs up to $M_{\rm EC} \approx 10^6$ $_{\sun}$ are often overlooked in surveys, much effort has been made to detect and to analyze stellar objects of $M \approx 10^7$ $_{\sun}$ since the discovery of UCDs in the Fornax Cluster by \cite{hilker99} and \cite{drinkwater00}."724. Therefore. the incompleteness for UCDs ts expected to be considerably lower.," Therefore, the incompleteness for UCDs is expected to be considerably lower."725 In conclusion. the rapidly increasing number of detectec ECs and UCDs associated with various types of galaxies i different environments offers à new perspective to the process of cluster formation and galaxy evolution.," In conclusion, the rapidly increasing number of detected ECs and UCDs associated with various types of galaxies in different environments offers a new perspective to the process of cluster formation and galaxy evolution."726 Since galaxy-galaxy mergers are anticipated to have been more common during early cosmological times it is expected that star-formation in cluster complexes has been a significant star-formation mode during this epoch., Since galaxy-galaxy mergers are anticipated to have been more common during early cosmological times it is expected that star-formation in cluster complexes has been a significant star-formation mode during this epoch.727 Our work provides a unification of the compact state of young star clusters and the extended size of ECs by allowing ECs to form from a CC of compact star clusters., Our work provides a unification of the compact state of young star clusters and the extended size of ECs by allowing ECs to form from a CC of compact star clusters.728 If the formation scenario suggested in this paper is correct. a statistical analysis of the observed ECs und UCDs with respect to the results of the models has the potential to shed light on the mass spectrum of the initial CCs and thereby on the interaction history of galaxies at this cosmologically important epoch.," If the formation scenario suggested in this paper is correct, a statistical analysis of the observed ECs und UCDs with respect to the results of the models has the potential to shed light on the mass spectrum of the initial CCs and thereby on the interaction history of galaxies at this cosmologically important epoch."729 Considerably larger. more homogeneous. and more complete datasets of ECs and UCDs and more detailed observations of CCs are. however. necessary to draw statistically significant conclusions on their origin.," Considerably larger, more homogeneous, and more complete datasets of ECs and UCDs and more detailed observations of CCs are, however, necessary to draw statistically significant conclusions on their origin."730pulsars in the first half of a new survey of the southeru Galactic plane usine the receutlv commissioned Parkes À 21-can unultibeam system.,pulsars in the first half of a new survey of the southern Galactic plane using the recently commissioned Parkes $\lambda$ 21-cm multibeam system.731 Their survey is some seven times more sensitive than the Clifton et al., Their survey is some seven times more sensitive than the Clifton et al.732 and Jolustou ct al., and Johnston et al.733" survevs, and the new discoveries already include severa ünarv pulsars (e.g. Lyne et al."," surveys, and the new discoveries already include several binary pulsars (e.g. Lyne et al."734 1999). as well as a large nuniber of very cistant. hieh cispersion measure. SOlCCS.," 1999), as well as a large number of very distant, high dispersion measure, sources."735 A preliminary account of the exciting results from the Parkes uultibeam survey (CamLB.o et al., A preliminary account of the exciting results from the Parkes multibeam survey (Camilo et al.736 1997) prompted us to utilise the laree collecting area of the 100nm Effelsherg radio telescope to perform à ucw search along the northern Galactic plane., 1997) prompted us to utilise the large collecting area of the 100–m Effelsberg radio telescope to perform a new search along the northern Galactic plane.737 In this paper we report on a σα] survey carried out during 1998 to test the feasibility of future observations with a wide-baudwilth search system currently under development., In this paper we report on a small survey carried out during 1998 to test the feasibility of future observations with a wide-bandwidth search system currently under development.738 This pilot search proved successful. discovering four new pulsus the first ever found with this telescope.," This pilot search proved successful, discovering four new pulsars — the first ever found with this telescope."739 Iu Sect., In Sect.740 2 we describe in sole detail the survey observations aud data reduction techniques., \ref{sec:obs} we describe in some detail the survey observations and data reduction techniques.741 In Sect., In Sect.742 5. we estimate the sensitivity of the survey., \ref{sec:sens} we estimate the sensitivity of the survey.743 The results are presented in Sect. L., The results are presented in Sect. \ref{sec:res}.744 These results. along with follow-up timing observations. are discussed in Sec. 5.," These results, along with follow-up timing observations, are discussed in Sec. \ref{sec:fup}."745 Finally. in Sect. 6..," Finally, in Sect. \ref{sec:conc},"746 we sunnmnarbsse the main conclusions frou this work and their iupliceatious for future pulsar search experiments at Efftelsbere., we summarise the main conclusions from this work and their implications for future pulsar search experiments at Effelsberg.747 All observations reported in this paper were carried out at a centre frequency of 1102. AUTz on a nuuber of separate sessions between 1998 June and 1999 April using the 1001) Effelsbere radio telescope operated by the Max-Plauck-Iustitut fürr Radioastrouomic., All observations reported in this paper were carried out at a centre frequency of 1402 MHz on a number of separate sessions between 1998 June and 1999 April using the 100–m Effelsberg radio telescope operated by the Max-Planck-Institut fürr Radioastronomie.748 Although LlOO-AMIz timing observations at Effelsbere are routinely made with typical baucwidths of 10 MIIz or more — (see e.g. Kramer et al., Although 1400-MHz timing observations at Effelsberg are routinely made with typical bandwidths of 40 MHz or more \nocite{kxl+98} (see e.g. Kramer et al.749 1998). the search hardware available to ΠΕ bandwidth of 16 MITz in cach ofthe two orthosonal. circular polarisation channels," 1998), the search hardware available to us has a maximum bandwidth of 16 MHz in each of the two orthogonal, circular polarisation channels."750 Nonetheless. the large forward gain of the telescope at 1100. ALTTz (1.5 K +). the relatively low svstem temperature of the receiver (35 IK) and long integration times eniploved in the survev (35 nün per pointing) means that the syste achieves a sensitivity which represcuts a threefold nuprovenient over that achiever bv Clifton et al. (," Nonetheless, the large forward gain of the telescope at 1400 MHz (1.5 K $^{-1}$ ), the relatively low system temperature of the receiver (35 K) and long integration times employed in the survey (35 min per pointing) means that the system achieves a sensitivity which represents a threefold improvement over that achieved by Clifton et al. ("7511992) durinug their survey.,1992) during their survey.752 The main aimi of the observations reported here was to test the feasibility of a larger search with a 100-MITz bandwidth svstem which is prescuthy being commissioned., The main aim of the observations reported here was to test the feasibility of a larger search with a 100-MHz bandwidth system which is presently being commissioned.753 Given the limited amount of telescope time available for this pilot project. we chose to restrict our searcli area to a 2 dee? patch of the Galactic plane defined by 28°«7«30° and |b|«0.57.," Given the limited amount of telescope time available for this pilot project, we chose to restrict our search area to a 2 $^2$ patch of the Galactic plane defined by $28^{\circ}<l<30^{\circ}$ and $|b|<0.5^{\circ}$."754 The rationale for this choice is sinyple this line-ofsielit is close to the Seutum spiral ari and as a result passes through one of the most pulsar-rich parts of the northern Calactic plane., The rationale for this choice is simple — this line-of-sight is close to the Scutum spiral arm and as a result passes through one of the most pulsar-rich parts of the northern Galactic plane.755 Iu addition. since this part of the sky is not visible from Arecibo. Effelsbere is preseutlv the largest radio telescope iu the world capable of surveviug it.," In addition, since this part of the sky is not visible from Arecibo, Effelsberg is presently the largest radio telescope in the world capable of surveying it."756" The survev region was divided up iuto a grid of 126 positions consisting of 9 strips of LL positions along lines of constant galactic latitude (b=0.0"",20.127.20.2L.40.367.£0.E87)."," The survey region was divided up into a grid of 126 positions consisting of 9 strips of 14 positions along lines of constant galactic latitude $b=0.0^{\circ}, \pm0.12^{\circ}, \pm0.24^{\circ}, \pm0.36^{\circ},757\pm0.48^{\circ}$ )."758 This choice of spacing ensured some overlap between the 3-dB width of the telescope can (9°)., This choice of spacing ensured some overlap between the 3-dB width of the telescope beam (9').759" The b=0.07 strip was centred ou |—29"".", The $b=0.0^{\circ}$ strip was centred on $l=29^{\circ}$.760 Beam ceutres on adjacent strips were altcruately offset bv half à beam width to ensure the most effücieut coverage on the sky., Beam centres on adjacent strips were alternately offset by half a beam width to ensure the most efficient coverage on the sky.761 At the start of cach observing run. woe carriccl out a S-anin observation of PSR D2011]238.," At the start of each observing run, we carried out a 5-min observation of PSR B2011+38."762 This relatively huninous 230-3108 pulsar has a dispersion ucasure of 239 cin? pc and is known not to be xone to significant iuteusity variations due to iuterstellar scintillation (Lorimer et al., This relatively luminous 230-ms pulsar has a dispersion measure of 239 $^{-3}$ pc and is known not to be prone to significant intensity variations due to interstellar scintillation (Lorimer et al.763 1995)., 1995).764 The fact that the search code detected this pulsar with consistently high sigual-to-lolse ratios (consistent with its L100-MITz £ux density 6.1405 inv: Loriuer et al., The fact that the search code detected this pulsar with consistently high signal-to-noise ratios (consistent with its 1400-MHz flux density — $6.4\pm0.5$ mJy; Lorimer et al.765 1995) eave us confidence hat the individual &lterbauk chauuels were functioning rormally. and that he nominal svsteni sensitivity was iue achieved.," 1995) gave us confidence that the individual filterbank channels were functioning normally, and that the nominal system sensitivity was being achieved."766 The search field is visible from Effeleberg for about 8.5 hr per dax., The search field is visible from Effelsberg for about 8.5 hr per day.767 Since cach evid position iu he field was observed for 35 mun. we typically observed up to 11 separate positions on the sky during a given transit.," Since each grid position in the field was observed for 35 min, we typically observed up to 14 separate positions on the sky during a given transit."768 Iu search node. the incoming signals of cach polarisation are fed iuto a xir of L LAMz filterbanks.," In search mode, the incoming signals of each polarisation are fed into a pair of $4\times4$ -MHz filterbanks."769 The outputs TOM the filterbauks are subsequently detected and digitised every 500 ps using 2-MITz voltage-to-frequeney converters. resulting iu an effective LO-bit quantisation of the signals.," The outputs from the filterbanks are subsequently detected and digitised every 500 $\mu$ s using 2-MHz voltage-to-frequency converters, resulting in an effective 10-bit quantisation of the signals."770 This is the fastest sustainable data rate using this svsteun., This is the fastest sustainable data rate using this system.771 Signals from the orthogonal polarisations were combined: to fori a total power time series for cach [-MITz frequency channel over the baud., Signals from the orthogonal polarisations were combined to form a total power time series for each 4-MHz frequency channel over the band.772 These four frequency chanuels are then passed to the standard Effelsbere Pulsar Observiug System (see ται 1995) which stored contiguous blocks of data to disk every 1021 samples (0.512 x)., These four frequency channels are then passed to the standard Effelsberg Pulsar Observing System (see \nocite{kra95} Kramer 1995) which stored contiguous blocks of data to disk every 1024 samples (0.512 s).773 The fom-channel search svstem used for this survev results in refreshingly low data rates compared to most other searches where the backends routinely sample 256 channels or more (sce c.g. Manchester ct al., The four-channel search system used for this survey results in refreshingly low data rates compared to most other searches where the backends routinely sample 256 channels or more \nocite{mld+96} (see e.g. Manchester et al.774 1996)., 1996).775 The nain advantage of such a simple svstei is that. as soon as cach 35-anin inteeration was complete. a preliminary analysis of the data could be carried out well within he tine that the telescope was observiug the next exid yosition.," The main advantage of such a simple system is that, as soon as each 35-min integration was complete, a preliminary analysis of the data could be carried out well within the time that the telescope was observing the next grid position."776 This quasi ou-the-fiv processing scheme allowed rapid re-observation of auv pulsar candidates found during he search., This quasi on-the-fly processing scheme allowed rapid re-observation of any pulsar candidates found during the search.777 The data analysis procedure was optimised to search Fourier spectra of cach 35-1in time series for dispersed seriodic signals., The data analysis procedure was optimised to search Fourier spectra of each 35-min time series for dispersed periodic signals.778 The software for this purpose was developed largely. frou scratch. taking advantage of ideas used to process our on-golug search of the Galactic centre from Effelshere (Usramer ct al.," The software for this purpose was developed largely from scratch, taking advantage of ideas used to process our on-going search of the Galactic centre from Effelsberg (Kramer et al."779 1996: Ixxiuuer et al., 1996; Kramer et al.780 2000). ax well as previous experience gained bv one," 2000), as well as previous experience gained by one"781provide here a short descriptions of the most relevant features of the code.,provide here a short descriptions of the most relevant features of the code.782 In order to maintain the general approach of the multiphase model by SHO3. we assume that stars with masses =40A. explode into SNII soon after their formation. thereby promptly releasing energy and metals.," In order to maintain the general approach of the multiphase model by SH03, we assume that stars with masses $>40\,M_\odot$ explode into SNII soon after their formation, thereby promptly releasing energy and metals."783 In contrast. we correctly accoun or the lifetime of stars having masses smaller than 40A7..," In contrast, we correctly account for the lifetime of stars having masses smaller than $40\,M_\odot$."784 The simulations that we will discuss here use the lifetimes providec by 2.. which have been shown to reproduce the abundance pattern in the Milky Way (2)..," The simulations that we will discuss here use the lifetimes provided by \cite{1989A&A...210..155M}, which have been shown to reproduce the abundance pattern in the Milky Way \citep{1997ApJ...477..765C}."785 Within the stochastic approach to star ormation (SHOQ3). each star particle is generated with a mass equa o one third of the mass of its parent gas particle.," Within the stochastic approach to star formation (SH03), each star particle is generated with a mass equal to one third of the mass of its parent gas particle."786 Therefore. each star particle is considered as a single stellar »opulation (SSP. hereafter). with its own mass. metallicity and redshift of formation.," Therefore, each star particle is considered as a single stellar population (SSP, hereafter), with its own mass, metallicity and redshift of formation."787 For each SSP we compute both the number of stars turning into SNII and Ta at each time-step and the number of stars ending their AGB phase., For each SSP we compute both the number of stars turning into SNII and Ia at each time-step and the number of stars ending their AGB phase.788 Then we calculate the amoun of energy and metals produced by each star particle in a given time interval. decreasing accordingly the mass of the particle.," Then we calculate the amount of energy and metals produced by each star particle in a given time interval, decreasing accordingly the mass of the particle."789 In this way. each star particle is characterized by both its initial mass. assigned at the time of its formation. and its final mass. which is updated during the evolution.," In this way, each star particle is characterized by both its initial mass, assigned at the time of its formation, and its final mass, which is updated during the evolution."790" Both SNIT and SNIa are assumed to release LO"" ergs each. while no energy output is associated to the mass loss from AGB stars."," Both SNII and SNIa are assumed to release $10^{51}$ ergs each, while no energy output is associated to the mass loss from AGB stars."791 The relative number of SNI and SNIa depends on the choice of the stellar initial mass function (IMF)., The relative number of SNII and SNIa depends on the choice of the stellar initial mass function (IMF).792" In the following. we will assume for the IMF the power-law shape dNidlogmxm""."," In the following, we will assume for the IMF the power–law shape $dN/d\log{m}\propto793m^{-x}$."794 Simulations will be run by assuming the Salpeter IMF with »=1.35 (2.Sa-IMFhereafter) and a heavy IMF with.»=0.95 (2.. TH-IMF hereafter).," Simulations will be run by assuming the Salpeter IMF with $x=1.35$ \citep[][Sa-IMF hereafter]{1955ApJ...121..161S} and a top--heavy IMF with $x=0.95$ \citealt{1987A&A...173...23A}, TH-IMF hereafter)."795 The SNIa are associated to binary systems whose components are in the Q.8-SA7. mass range (2).. while SNII arise from stars with mass >S A/..," The SNIa are associated to binary systems whose components are in the $8\,M_\odot$ mass range \citep{1983A&A...118..217G}, while SNII arise from stars with mass $>8\,M_\odot$ ."796 In the following. we will assume that 10 per cent of stars in the 0.8—5AZ. mass range belongs to binary systems. which then produces SNIa.," In the following, we will assume that 10 per cent of stars in the $8\,M_\odot$ mass range belongs to binary systems, which then produces SNIa."797We use the analytical fitting formulas for stellar yields of SNIu. SNII and PNe provided by ?.. and based on the original nucleosynthesis computations of 2.. using their W7 model. ? and ?..,"We use the analytical fitting formulas for stellar yields of SNIa, SNII and PNe provided by \cite{2001MNRAS.322..800R}, and based on the original nucleosynthesis computations of \cite{1997NuPhA.621..467N}, , using their W7 model, \cite{1995ApJS..101..181W} and \cite{1981A&A....94..175R}."798 The formulation for the SNIa rate has been calculated as in 2.., The formulation for the SNIa rate has been calculated as in \cite{2001ApJ...558..351M}.799 In the simulations that we present. besides H and He. we have followed Fe. O. C. Si. Mg. S. Once produced by a star particle. metals are spread over the same number of neighbours. 64. used for the SPH computations. also using the same kernel.," In the simulations that we present, besides H and He, we have followed Fe, O, C, Si, Mg, S. Once produced by a star particle, metals are spread over the same number of neighbours, 64, used for the SPH computations, also using the same kernel."800 We normalize the IMFs in the mass range 0.1—100 δι., We normalize the IMFs in the mass range 0.1–100 $M_\odot$.801 Owing to the uncertainty in modelling yields for very massive stars. we take yields to be independent of mass above 40A/..," Owing to the uncertainty in modelling yields for very massive stars, we take yields to be independent of mass above $40\,M_\odot$."802 While any uncertainty in the yields of such massive stars has a negligible effect for the Salpeter IMF. their accurate description (e.g.2?) is required when using a top-heavier IMF.," While any uncertainty in the yields of such massive stars has a negligible effect for the Salpeter IMF, their accurate description \citep[e.g.,][]{1996ApJ...460..408T,2002ApJ...567..532H}803 is required when using a top–heavier IMF."804" Our prescription to account for stellar evolution in the simulations implies a substantial change of the multiphase ""effective model” by SHO3. which we have suitably modified to account for(4) the contribution of the energy reservoir provided by the SN which are treated outside the TRA. and the metal-dependence of the cooling function. that we introduce using the tables from ?.."," Our prescription to account for stellar evolution in the simulations implies a substantial change of the multiphase “effective model” by SH03, which we have suitably modified to account for the contribution of the energy reservoir provided by the SN which are treated outside the IRA, and the metal–dependence of the cooling function, that we introduce using the tables from \cite{1993ApJS...88..253S}."805" The resulting density threshold. for a gus particle to become multiphase. thereby being eligible to undergo star formation. is fixed tony=OL em"" at zero metallicity."," The resulting density threshold for a gas particle to become multiphase, thereby being eligible to undergo star formation, is fixed to $n_H=0.1$ $^{-3}$ at zero metallicity."806 According to eq.(23) of SHO3. this. threshold is inversely proportional to the cooling function.," According to eq.(23) of SH03, this threshold is inversely proportional to the cooling function."807 Since the latter depends on metallicity. we take self-consistently a metallicity—dependent star-formation threshold for gas having a non-zero metallicity.," Since the latter depends on metallicity, we take self–consistently a metallicity--dependent star–formation threshold for gas having a non-zero metallicity."808 SH03 also provided a phenomenological description for galactic winds. which are triggered by SN energy release and whose strength is regulated by two parameters.," SH03 also provided a phenomenological description for galactic winds, which are triggered by SN energy release and whose strength is regulated by two parameters."809" The first one gives the wind mass loading according to the relation. My=DRM where AJ, is the star formation rate."," The first one gives the wind mass loading according to the relation, $\dot M_W=\eta \dot810M_*$, where $\dot M_*$ is the star formation rate."811 Following SHO3. we assume η=3.," Following SH03, we assume $\eta=3$."812 The second parameter is the wind velocity. ος.," The second parameter is the wind velocity, $v_w$."813 For the runs based on the Salpeter IMF. we always use 7.=500κι," For the runs based on the Salpeter IMF, we always use $v_w=500\vel$."814 For the above values of jj and ον. all the energy from SNIT is converted in kinetic energy. as in the original SHO3 paper.," For the above values of $\eta$ and $v_w$, all the energy from SNII is converted in kinetic energy, as in the original SH03 paper."815 ? made a study of the star formation history predicted by hydrodinamical simulations which include galactic winds with a similar velocity., \cite{2003MNRAS.339..312S} made a study of the star formation history predicted by hydrodinamical simulations which include galactic winds with a similar velocity.816 They concluded that the resulting star fraction at >=OQ and high—: star formation history are comparable to the observed ones., They concluded that the resulting star fraction at $z=0$ and $z$ star formation history are comparable to the observed ones.817 In order to verify the effect of galactic ejecta. the ¢676 and g51 regions are also simulated with the Salpeter IMF. but setting to zero the wind velocity (Sa-NW runs).," In order to verify the effect of galactic ejecta, the g676 and g51 regions are also simulated with the Salpeter IMF, but setting to zero the wind velocity (Sa-NW runs)."818" As for the runs based on the top— IMF. we will also use ey=500kms+ for all clusters. with the exception of 2676 and g51 regions. for which we also use e,=1000+ CTH-SW runs)."," As for the runs based on the top--heavy IMF, we will also use $v_W=500\vel$ for all clusters, with the exception of g676 and g51 regions, for which we also use $v_w=1000\vel$ (TH-SW runs)."819 In this case. the two wind speeds correspond kms.to an energy budget of about 0.4 and 1.4 times the energy provided by SNII.," In this case, the two wind speeds correspond to an energy budget of about 0.4 and 1.4 times the energy provided by SNII."820 An efficiency larger than unity can be justified on the ground of the large uncertainties on the actua energy released by SNII explosions., An efficiency larger than unity can be justified on the ground of the large uncertainties on the actual energy released by SNII explosions.821 In this perspective. we also take the value of the wind velocity as a confidence value.," In this perspective, we also take the value of the wind velocity as a confidence value."822 A wind velocity my=1000kms+ is intended to represent an extreme feedback case. so that comparing the results with the two values of wind velocity allows us to check the effect of a stronger feedback on the final properties of the galaxy population.," A wind velocity $v_W=1000\vel$ is intended to represent an extreme feedback case, so that comparing the results with the two values of wind velocity allows us to check the effect of a stronger feedback on the final properties of the galaxy population."823 We summarize in Table 2. the IMFs and feedback used in our simulations., We summarize in Table \ref{t:sim} the IMFs and feedback used in our simulations.824 As a first step. we identify galaxies from the distribution of star particles by applying the SKID (2)..," As a first step, we identify galaxies from the distribution of star particles by applying the SKID \citep{2001PhDT........21S}."825 We provide a short description of how we applied this algorithm. while a more detailed discussion and presentation of tests is provided elsewhere (Murante et al.," We provide a short description of how we applied this algorithm, while a more detailed discussion and presentation of tests is provided elsewhere (Murante et al."826 2006. in preparation: see also ?2)).," 2006, in preparation; see also \citealt{2006MNRAS.367.1641B}) )."827 An overall density field is computed by using the distribution of all the particle species. by using a SPH spline—kernel.," An overall density field is computed by using the distribution of all the particle species, by using a SPH spline--kernel."828 The star particles are then moved along the gradient of the density field in steps of 7/2. where we assume Τ—3ci4.," The star particles are then moved along the gradient of the density field in steps of $\tau/2$, where we assume $\tau\simeq 3 \epsilon_{\rm Pl}$."829 When a particle begins to oscillate inside a sphere of radius 7/2. it is stopped.," When a particle begins to oscillate inside a sphere of radius $\tau/2$, it is stopped."830 Once all particles have been moved. they are grouped using a friends-of-friends (FOF) algorithm. with linking length 7/2. applied to the new particle positions.," Once all particles have been moved, they are grouped using a friends-of-friends (FOF) algorithm, with linking length $\tau/2$, applied to the new particle positions."831 The binding energy of each group identifiedin this way is then used to remove from the group all star particles which are recognized as unbound., The binding energy of each group identifiedin this way is then used to remove from the group all star particles which are recognized as unbound.832 All particles in a sphere of radius 7. centered on the centerof mass of the group. are used to compute such a gravitational binding," All particles in a sphere of radius $\tau$ , centered on the centerof mass of the group, are used to compute such a gravitational binding"833Ginga. aud recentBeppoSAN X-ray observatious of Sevtert 1. ealaxies strouely support a eeneralobservational framework that meludes: ionized absorption by a warn absorber (WA). a steep intrinsic power-law coutinmun. a Compton reflection hunip and associated neutral Felva line (narrow aud/or broad). aud a ligh-cnerey cutoff (Lear).,", and recent X-ray observations of Seyfert 1 galaxies strongly support a general framework that includes: ionized absorption by a warm absorber (WA), a steep intrinsic power-law continuum, a Compton reflection hump and associated neutral ${\alpha}$ line (narrow and/or broad), and a high-energy cutoff $E_{\rm cutoff}$ )."834 These observational results have raised the following questious:, These observational results have raised the following questions:835 Finally. Fig.," Finally, Fig."836 4. shows the 250um profiles perpendicular to the major axis. for three cuts passing through the galaxy centre and the two secondary peaks at about I’ from the centre on the NE and SW.," \ref{figver} shows the $\mu$ m profiles perpendicular to the major axis, for three cuts passing through the galaxy centre and the two secondary peaks at about $\arcmin$ from the centre on the NE and SW."837 Data were averaged over two beams perpendicular to the cut. and a residual background was subtracted from regions beyond the emission.," Data were averaged over two beams perpendicular to the cut, and a residual background was subtracted from regions beyond the emission."838 Model profiles were scaled to match the data on the major axis., Model profiles were scaled to match the data on the major axis.839 Emission above the plane is almost entirely dominated by the PSF Ay rings picking up light from the major axis (because the adopted dust vertical scalelength. 0.2 kpe or 4”. is much smaller than the main beam).," Emission above the plane is almost entirely dominated by the PSF Airy rings picking up light from the major axis (because the adopted dust vertical scalelength, 0.2 kpc or $\arcsec$, is much smaller than the main beam)."840 Results are similar in the other SPIRE bands., Results are similar in the other SPIRE bands.841 Excess emission is only seen on the NW side of the minor axis profile. in coincidence with the background X-ray source CXO J022224.44-422138. which contributes to most of the excess: however. part of it could come from another background source aligned with the minor axis or from a filamentary structure protruding from the galactic centre.," Excess emission is only seen on the NW side of the minor axis profile, in coincidence with the background X-ray source CXO J022224.4+422138, which contributes to most of the excess; however, part of it could come from another background source aligned with the minor axis or from a filamentary structure protruding from the galactic centre."842 Excluding this feature. SPIRE images do not show evidence of any diffuse halo emission.," Excluding this feature, SPIRE images do not show evidence of any diffuse halo emission."843 This is at odds with the previous Spitzer observations of ?.. which detected MIR. emission extending up to 5 kpe from the plane.," This is at odds with the previous Spitzer observations of \citet{BurgdorfApJ2007}, which detected MIR emission extending up to 5 kpc from the plane."844 At 22um. a surface brightness of 0.1-0.2 MJy/sr is found at 1-144 along the minor axis.," At $\mu$ m, a surface brightness of 0.1-0.2 MJy/sr is found at 4 along the minor axis."845 Instead. at 250.m. residual emission from the halo ts less than | MJy/sr.," Instead, at $\mu$ m, residual emission from the halo is less than 1 MJy/sr."846 Such MIR/submm flux density ratios can be produced by typical Milky Way dust. if the interstellar radiation field is 10x higher than the local (see. e.g.. Fig.," Such MIR/submm flux density ratios can be produced by typical Milky Way dust, if the interstellar radiation field is $\times$ higher than the local (see, e.g., Fig."847 13 in 2))., 13 in \citealt{DraineApJ2007b}) ).848 These heating conditions might be compatible with the hypothesis that the detected MIR radiation come from dust surrounding halo AGB stars (?).., These heating conditions might be compatible with the hypothesis that the detected MIR radiation come from dust surrounding halo AGB stars \citep{BurgdorfApJ2007}.849 We have analysed SPIRE observations of the edge-on spiral galaxy NGC 891 and examined both their contribution to the galaxy’s SED and the galaxy’s morphological properties at these wavelengths., We have analysed SPIRE observations of the edge-on spiral galaxy NGC 891 and examined both their contribution to the galaxy's SED and the galaxy's morphological properties at these wavelengths.850 Emission within the central 44 broadly follows a radial exponential distribution with a scalelength similar to those derived for the stars in the surface brightness fits of ?.., Emission within the central $\farcm$ 4 broadly follows a radial exponential distribution with a scalelength similar to those derived for the stars in the surface brightness fits of \citet{XilourisSub1998}.851 Emission above the galactic plane can be explained with radiation from the galactic plane picked up by the PSF wings., Emission above the galactic plane can be explained with radiation from the galactic plane picked up by the PSF wings.852 The distribution of dust along the vertical direction ts thus thinner than the SPIRE PSFs. and compatible with the vertical scalelength derived by ?..," The distribution of dust along the vertical direction is thus thinner than the SPIRE PSFs, and compatible with the vertical scalelength derived by \citet{XilourisSub1998}."853 Beyond 4/4 from the centre. the radial profile shows a break and becomes steeper.," Beyond $\farcm$ 4 from the centre, the radial profile shows a break and becomes steeper."854 On the NE side. the submm disk comes to an end within the optical size of the galaxy.," On the NE side, the submm disk comes to an end within the optical size of the galaxy."855 Instead. on the SW side the profile decline is interrupted at about the optical radius.," Instead, on the SW side the profile decline is interrupted at about the optical radius,"856"the ionized In principle then, the ionization should be able to significantly influence the evolution of the protocluster's gas reserves.","the ionized In principle then, the ionization should be able to significantly influence the evolution of the protocluster's gas reserves."857" In Figure 12, we plot the positions of the gas whose density is too large to be affected by the HII."," In Figure \ref{fig:toodense}, we plot the positions of the gas whose density is too large to be affected by the HII."858" We neglect here the ram-pressure of the cold gas against which the HII must also contend, since the gas is infalling towards the ionizing sources."," We neglect here the ram–pressure of the cold gas against which the HII must also contend, since the gas is infalling towards the ionizing sources."859 The dense gas clearly (and not surprisingly) correlates strongly with the positions of the clusters in the simulation — the stars are to be found embedded in the densest gas., The dense gas clearly (and not surprisingly) correlates strongly with the positions of the clusters in the simulation – the stars are to be found embedded in the densest gas.860" In order to significantly affect the evolution of the cluster, photoionization would need to clear this material away."," In order to significantly affect the evolution of the cluster, photoionization would need to clear this material away."861" Since it can do so only very slowly, as shown in Figures 7 and 8,, the hot gas instead leaks into the low-density Classical models (e.g??7?) indicate that ionization—driven champagne flows should be an efficient mechanism for disrupting molecular clouds and this is very likely to be true in many systems, such as Orion (?),, where O-stars are located near the edges of clouds."," Since it can do so only very slowly, as shown in Figures \ref{fig:aniso1} and \ref{fig:aniso2}, the hot gas instead leaks into the low–density Classical models \citep[e.g][]{1979ApJ...233...85B,1987sbge.proc..467L,2002ApJ...566..302M} indicate that ionization--driven champagne flows should be an efficient mechanism for disrupting molecular clouds and this is very likely to be true in many systems, such as Orion \citep{2005ApJ...627..813H}, where O–stars are located near the edges of clouds."862" However, in cases where O-stars are embedded in highly inhomogeneous and often filamentary gas deep in cloud interiors, photoionization may not be so effective."," However, in cases where O–stars are embedded in highly inhomogeneous and often filamentary gas deep in cloud interiors, photoionization may not be so effective."863" In this work and that of several other authors (e.g???7),, embedded O-stars struggle to disrupt their natal clusters by photoionization."," In this work and that of several other authors \citep[e.g][]{1989A&A...216..207Y, 2005MNRAS.358..291D,2010ApJ...711.1017P,2010ApJ...719..831P}, embedded O–stars struggle to disrupt their natal clusters by photoionization."864" This does not imply that this must always be the case, but it does suggest that gas expulsion by ionizing radiation is more difficult than implied by earlier models."," This does not imply that this must always be the case, but it does suggest that gas expulsion by ionizing radiation is more difficult than implied by earlier models."865" This difference emerges principally because strong accretion flows onto the ionizing sources suppress the growth of HII regions, a process which was previously often Stellar winds may assist the UV radiation (e.g??7),, although we found in ? that cluster disruption by winds alone may be a rather slow process."," This difference emerges principally because strong accretion flows onto the ionizing sources suppress the growth of HII regions, a process which was previously often Stellar winds may assist the UV radiation \citep[e.g][]{1997A&A...326.1195C,2001PASP..113..677C,2003ApJ...594..888F}, although we found in \cite{2008MNRAS.391....2D} that cluster disruption by winds alone may be a rather slow process."866" The oldest O-stars in our simulation will soon explode as supernovae and it is possible that the mechanical and energetic feedback from the few hundred O-stars that reside in our clusters would succeed where photoionization has failed. ?,,"," The oldest O–stars in our simulation will soon explode as supernovae and it is possible that the mechanical and energetic feedback from the few hundred O–stars that reside in our clusters would succeed where photoionization has failed. \cite{1985A&A...145...70T},"867" for example, concluded that a single supernova could disrupt ~10* of molecular cloud material."," for example, concluded that a single supernova could disrupt $\sim10^{4}$ of molecular cloud material."868" The gravitational binding energy of our model system is a fewx10*! erg and it has formed the equivalent of a few hundred O-stars by an age of 6Myr, so that supernovae should in principle have no trouble unbinding it."," The gravitational binding energy of our model system is a $\times10^{51}$ erg and it has formed the equivalent of a few hundred O–stars by an age of 6Myr, so that supernovae should in principle have no trouble unbinding it."869" Alternatively, as suggested and investigated by ?,, molecular clouds and their embedded clusters (or cluster complexes) may not be bound in the first place."," Alternatively, as suggested and investigated by \cite{2005MNRAS.359..809C}, molecular clouds and their embedded clusters (or cluster complexes) may not be bound in the first place."870" JED acknowledges support from a Marie Curie fellowship as part of the European Commission FP6 Research Training Network ‘Constellation’ under contract MRTN-CT-2006-035890, and from the Institutional Research Plan AV0Z10030501 of the Academy of Sciences of the Czech Republic and project LC06014—Centre for Theoretical Astrophysics of the Ministry of Education, Youth and Sports of the Czech Republic."," JED acknowledges support from a Marie Curie fellowship as part of the European Commission FP6 Research Training Network `Constellation' under contract MRTN--CT--2006--035890, and from the Institutional Research Plan AV0Z10030501 of the Academy of Sciences of the Czech Republic and project LC06014–Centre for Theoretical Astrophysics of the Ministry of Education, Youth and Sports of the Czech Republic."871been completed such that they would include the Blancllord-Znajek mechanism.,been completed such that they would include the Blandford-Znajek mechanism.872 The results presented. in this paper are dependent on our assumptions that a BlI-disc magnetic connection exists., The results presented in this paper are dependent on our assumptions that a BH-disc magnetic connection exists.873 Closed. magnetic field lines in the BIL ergosphere may. be produced by a current ring in the vicinity of the DII., Closed magnetic field lines in the BH ergosphere may be produced by a current ring in the vicinity of the BH.874 Models for the magnetic connection where a poloidal magnetic field is generated. by a single electric current Llowing in the BL equatorial plane or at the inner edge of the accretion disc were proposed. for instance. hy ο and ?..," Models for the magnetic connection where a poloidal magnetic field is generated by a single electric current flowing in the BH equatorial plane or at the inner edge of the accretion disc were proposed, for instance, by \citet{li02R} and \citet{wang}."875 The kev parameters of the proposed model are. however. the 111 mass. the 111 spin ancl the mass aceretion rate.," The key parameters of the proposed model are, however, the BH mass, the BH spin and the mass accretion rate."876 One indirect way to test whether this mechanism operates in reality is to study the relation between the observed. radio Bux densitv from AGN (ο. from. Lat-spectrum core sources) and their mass accretion rates in order to fit the model prediction with respect to the relation tween the power of the jet and. the mass accretion rate (sce Fie. 4)).," One indirect way to test whether this mechanism operates in reality is to study the relation between the observed radio flux density from AGN (e.g., from flat-spectrum core sources) and their mass accretion rates in order to fit the model prediction with respect to the relation between the power of the jet and the mass accretion rate (see Fig. \ref{PjetsVsMdot}) )."877 That relation shows that the power of the jet does not depend linearly on the mass accretion rate all the way down to very low acerction rates. so that there can oe sources with relatively strong jet power but low mass accretion rate.," That relation shows that the power of the jet does not depend linearly on the mass accretion rate all the way down to very low accretion rates, so that there can be sources with relatively strong jet power but low mass accretion rate."878 In this case the jet power is mainly dependent on the DII parameters. such as the mass and the spin of the DII.," In this case the jet power is mainly dependent on the BH parameters, such as the mass and the spin of the BH."879 The first step is to find a relation between the power of the jets and the observed. radio flux density. Z5.," The first step is to find a relation between the power of the jets and the observed radio flux density, $F_{\rm obs}$."880 For a conical (where the magnetic field along the jet scales with he distance.jet B~2+t. and the electron number density. in he jet scales as C~2 2) from a Hat-spectrum core source. we found the dependence: Pas.71277DEOAE.1719. where D is the distance to the AGN and Al is the DII mass.," For a conical jet (where the magnetic field along the jet scales with the distance, $B \sim z^{-1}$, and the electron number density in the jet scales as $C' \sim z^{-2}$ ) from a flat-spectrum core source, we found the dependence: $P_{\rm jets} \sim F_{\rm obs}^{6/5} D^{12/5} M^{-7/10}$, where $D$ is the distance to the AGN and $M$ is the BH mass."881 A ull description of its derivation is the subject to a paper in preparation., A full description of its derivation is the subject to a paper in preparation.882 The second step is to produce a complete sample of ACN with known jet parameters. as the Doppler actor. and whose mass acerction rate can be constrained by observational data. and then to fit the model prediction.," The second step is to produce a complete sample of AGN with known jet parameters, as the Doppler factor, and whose mass accretion rate can be constrained by observational data, and then to fit the model prediction."883 Furthermore. if AGN were in the Povnting flux limit disregarding the mass accretion rate. then one might expect that the power in the jet [rom a large sample runs into a lower limit. which is given by the minimum DPovnting flux.," Furthermore, if AGN were in the Poynting flux limit disregarding the mass accretion rate, then one might expect that the power in the jet from a large sample runs into a lower limit, which is given by the minimum Poynting flux."884 The model presented. here can also be extended: to microquasars assuming that physical quantities (e.g. DII magnetic field. jet power. ete.)," The model presented here can also be extended to microquasars assuming that physical quantities (e.g., BH magnetic field, jet power, etc.)"885 scale with the BIL mass., scale with the BH mass.886 Although numerical simulations of jet. formation fron the crgosphere of a rapidly-spinning BIL with closed magnetic field lines that connect the DII to the cise inside the ergosphere has not been performed vet. this can be one of the challenges to be faced by numerical relativists.," Although numerical simulations of jet formation from the ergosphere of a rapidly-spinning BH with closed magnetic field lines that connect the BH to the disc inside the ergosphere has not been performed yet, this can be one of the challenges to be faced by numerical relativists."887 The author would like to thank Peter L. Diermann., The author would like to thank Peter L. Biermann.888 This research was supported through a stipend. from the International Max Planck Research School. (IMPIIS) for Astronomy and Astrophysics at. the Universities of Bonn ane Ixólln., This research was supported through a stipend from the International Max Planck Research School (IMPRS) for Astronomy and Astrophysics at the Universities of Bonn and Kölln.889 The author appreciates the support from MPLUR during the last phase of this work., The author appreciates the support from MPIfR during the last phase of this work.890 llere. we present the derivation of the angular monentum. conservation law for the matter that [lows in the veeretion disc. (eq. 10)).," Here, we present the derivation of the angular momentum conservation law for the matter that flows in the accretion disc (Eq. \ref{eq:angmom}) )."891. This derivation is based. on Ye general-relativistic angular-momoentum-conservation [aw ja describes the structure of a geometricallv thin acerction isc) (?).., This derivation is based on the general-relativistic angular-momentum-conservation law that describes the structure of a geometrically thin accretion disc \citep{pt}.892 απο on this conservation law. ? clerivec 16 conservation law that includes the Bll-clise magnetic 'onnection. and ?. derived the conservation law that includes rw jet.," Based on this conservation law, \citet{li02} derived the conservation law that includes the BH-disc magnetic connection, and \citet{db} derived the conservation law that includes the jet."893 In a slightly dillerent manner. our result is obtainec when both the Bll-dise magnetic connection and the je ormation are considered.," In a slightly different manner, our result is obtained when both the BH-disc magnetic connection and the jet formation are considered."894 The angular-momentum conservation law of the matter in à standard. thin accretion disc is given by eq. (," The angular-momentum conservation law of the matter in a standard, thin accretion disc is given by eq. ("89523) of ?7.. ?..,23) of \citet{pt}. \citet{pt}.896 11. 14)) AT , \ref{torque} \ref{eq:li}) \ref{23} 897the redder pixels in region (b) described above seen to be caused by thick dust lances ruunius from tle NW to SE in the galaxy.,the redder pixels in region (b) described above seem to be caused by thick dust lanes running from the NW to SE in the galaxy.898 It is hard to see whether regious (a) and (0) are caused by a dust lane or not. but regions (a) seclus to be distributed around bluer chirps in3.," It is hard to see whether regions (a) and (c) are caused by a dust lane or not, but regions (a) seems to be distributed around bluer clumps in."899. The localized regions (ο) do not appear to be unrelated backgrouud or foreground objects., The localized regions (c) do not appear to be unrelated background or foreground objects.900 Such objects usually have colors sufficicutly differeut from those of the eeuuine ealaxy pixels. that they would appear as a distiuct aud separate branch or grouping of pixels iu pCMDs and pCCDs.," Such objects usually have colors sufficiently different from those of the genuine galaxy pixels, that they would appear as a distinct and separate branch or grouping of pixels in pCMDs and pCCDs."901 Since we found uo such feature in the examined pCAIDs and pCCDs for 00959 (0... 1)). we treat these regions as parts of the galaxy.," Since we found no such feature in the examined pCMDs and pCCDs for 0959 (e.g., ), we treat these regions as parts of the galaxy."902 To studv the effect of our pixel-based extinction correction dn a lore quantitative way. we examined pCMDs and pCCDs using various combinations of iuages from FUV through ffor significant features or eroupines iu the pixeldistribution.," To study the effect of our pixel-based extinction correction in a more quantitative way, we examined pCMDs and pCCDs using various combinations of images from FUV through for significant features or groupings in the pixel-distribution."903 A «distinct grouping of pixels in these diagrams should indicate that those pixels are douinated bv the same or a similar nüx of stellar populations., A distinct grouping of pixels in these diagrams should indicate that those pixels are dominated by the same or a similar mix of stellar populations.904 Aduone differeut diagrams examined (not shown here). a pCCD of vvorsus ccolor extinction correction 1((0))) was selected. because it clearly shows distinct tracks aud eroupiues of pixels.," Among different diagrams examined (not shown here), a pCCD of versus color extinction correction )) was selected, because it clearly shows distinct tracks and groupings of pixels."905 Iu. particular. distiuct red aud blue sequences ean be identified.," In particular, distinct red and blue sequences can be identified."906 The uncertainties iu the colors at the onth and ΤΟΤΕ percentiles are ~ 00.01 and ~ nunag inP). and 00.05 and ~O0.08munag inpu). respectively.," The uncertainties in the colors at the $^{\rm th}$ and $^{\rm th}$ percentiles are $\sim$ 0.01 and $\sim$ mag in, and $\sim$ 0.05 and $\sim$ mag in, respectively."907=~ Since the gap of ~ 1inag at the bluer cud of ccolor between the two sequences in .((0)) is much larger than the photometric uncertainties in the colors. this separation of pixels is not a rancom effect.," Since the gap of $\sim$ mag at the bluer end of color between the two sequences in ) is much larger than the photometric uncertainties in the colors, this separation of pixels is not a random effect."908 The sequences are also not the result of differences in residual reddenius. as can be secu by comparing the directions of the sequences with that of the reddenimg vector drawn in l.," The sequences are also not the result of differences in residual reddening, as can be seen by comparing the directions of the sequences with that of the reddening vector drawn in )."909((e)) The red and blue sequences are connected in the = πας color range by pixels with intermediate colors., The red and blue sequences are connected in the $\lesssim$ $\lesssim$ mag color range by pixels with intermediate colors.910ος At colors redder than cm uunae. the pixels again somewhat separate into sequences with redder aud bluer ccolors.," At colors redder than $\simeq$ mag, the pixels again somewhat separate into sequences with redder and bluer colors."911 As can be seen iu various SED models Alaraston2005:I&otullaetal. 2009).. the ccolor is very seusitive to age for the vouugest stellar populations.," As can be seen in various SED models \citep[e.g.,][]{bruzual03, anders03,912maraston05, kotulla09}, the color is very sensitive to age for the youngest stellar populations."913 Since most of the FUV. flux is euütted by voung. massive OB-stars. combining a aud an optical broad-baud filter provides strong age coustraints (e.g.Navirajetal.2007).," Since most of the FUV flux is emitted by young, massive OB-stars, combining a and an optical broad-band filter provides strong age constraints \citep[e.g.,][]{kaviraj07}."914. The ccolor was selected eupircallv. while examining different combinations of colors.," The color was selected empirically, while examining different combinations of colors."915 The IRAC bbaudpass is conunonlv usec as a stellar παν» distribution tracer (0.9.Willueretal.2001).. because it is associated more with the distribution of redder and older stars.," The IRAC bandpass is commonly used as a stellar mass distribution tracer \citep[e.g.,][]{ willner04}, because it is associated more with the distribution of redder and older stars."916 Since the optical B-haud is ecucrally sensitive to vounger stars. the ccolor can be used to distinguish mixtures of stellar populations with aud without significaut recent lnass star formation.," Since the optical -band is generally sensitive to younger stars, the color can be used to distinguish mixtures of stellar populations with and without significant recent high-mass star formation."917 The combination of aand ccolors in .((6)). therefore. provides a powerful diagnostic of the recent star formation history averaged over a pixel.," The combination of and colors in ), therefore, provides a powerful diagnostic of the recent star formation history averaged over a pixel."918 Based on the distinct red and blue sequences. as well as transition regions identified iu d((a)). we separated pixels into six differcut groups (color-coded in £(0))).," Based on the distinct red and blue sequences, as well as transition regions identified in ), we separated pixels into six different groups (color-coded in ))."919 The group boundaries (selection criteria) are summarized in Table 1., The group boundaries (selection criteria) are summarized in Table 1.920 The exact location of the bouudaries between some of these pixel groups is solmewhat arbitrary. but is motivated by the following considerations.," The exact location of the boundaries between some of these pixel groups is somewhat arbitrary, but is motivated by the following considerations."921 The Croup I pixels (color-coded purple) reside ou the blue part of the blue sequence. aud have blue colors iu both aand)).," The Group I pixels (color-coded purple) reside on the blue part of the blue sequence, and have blue colors in both and."922. The dominant stellar populations are very young and massive OB-stars. while older stellar populations contribute little to the total flux in these pixels.," The dominant stellar populations are very young and massive OB-stars, while older stellar populations contribute little to the total flux in these pixels."923 The aand ccolors become redder as these voung stellar populations age and the wmuhber of remaining OB-stars decreases., The and colors become redder as these young stellar populations age and the number of remaining OB-stars decreases.924 On the other extreme. the Grows VI pixels coded red) have red colors in both aandο.," On the other extreme, the Group VI pixels (color-coded red) have red colors in both and."925" These pixels are therefore dominated by old. quiescent. ""red-aud-dead stellar populations aud do not contain a detectable fraction of voung. massive stars."," These pixels are therefore dominated by old, quiescent, “red-and-dead” stellar populations and do not contain a detectable fraction of young, massive stars."926 The Group II pixels (color-coded blue) have the same blue ccolor range as Croup I but redder ccolors.," The Group II pixels (color-coded blue) have the same blue color range as Group I, but redder colors."927 The blue ccolor implies the presence of OB-stars., The blue color implies the presence of OB-stars.928 The redder ccolor indicates that Croup ID pixels lave a non-ueeheible contribution to the total light from older Guuderlving or superposed aloug the lne-ofsight} stellar populations., The redder color indicates that Group II pixels have a non-negligible contribution to the total light from older (underlying or superposed along the line-of-sight) stellar populations.929 Group V pixels (color-coded orange) cover the sanie. red. ccolor range as Croup VI. indicating that they too lack a detectable fraction of vouug. massive OB-stars.," Group V pixels (color-coded orange) cover the same, red, color range as Group VI, indicating that they too lack a detectable fraction of young, massive OB-stars."930 Yet. heir blucr ccolor sugeests that the fiux in these pixels is dominated x hielt from intermediate age stellar populations.," Yet, their bluer color suggests that the flux in these pixels is dominated by light from intermediate age stellar populations."931 The Caoup TT and Croup IV pixels (color-coded ereen and vellow. respectively) are located between hese extreme cases.," The Group III and Group IV pixels (color-coded green and yellow, respectively) are located between these extreme cases."932 The lieht iu these pixels is likely dominated by stellar populations m transition between he extreme eroups along either blue or red sequences. or between the blue aud the red sequence. or represents a uixture of stellar populations (fom unresolved adjaceut regions or regions superposed along the lue-ofsight)," The light in these pixels is likely dominated by stellar populations in transition between the extreme groups along either blue or red sequences, or between the blue and the red sequence, or represents a mixture of stellar populations (from unresolved adjacent regions or regions superposed along the line-of-sight)"933redshift as isophotal magnitudes.,redshift as isophotal magnitudes.934 Thus the sample is selected rom pscuclo-total limits., Thus the sample is selected from pseudo-total limits.935 For deep isophotes. the volume correction has virtually no surface brightness dependence. as shown in Fig 5.. so he number density can be caleulatec at. the bright. end rivially and is only underestimated at the faint end where some galaxies have too a low surface brightness to get into the sample.," For deep isophotes, the volume correction has virtually no surface brightness dependence, as shown in Fig \ref{fig:lfs}, so the number density can be calculated at the bright end trivially and is only underestimated at the faint end where some galaxies have too a low surface brightness to get into the sample."936 Using either of the techniques outlined in he previous paragraph (Cross et al., Using either of the techniques outlined in the previous paragraph (Cross et al.937 2001. Blanton ct al.," 2001, Blanton et al."938 2001) should. give accurate values of A.Sloeh and ó*. §.7540.05 mag. (2.0220.02)1022 5 for 24dEGIBRS (Cross et al.," 2001) should give accurate values of $M^*_{b_j}-5\log\,h$ and $\phi^*$, $-19.75 \pm 0.05$ mag, $(2.02 939\pm 0.02)\times10^{-2}h^3$ $^{-3}$ for 2dFGRS (Cross et al."940 2001) and. 19.70c0.04 mae. (2.05+0.12)10ΗΝ * for SDSS (Blanton et al.," 2001) and $-19.70 \pm 0.04$ mag, $(2.05 \pm 0.12)\times10^{-2}h^3$ $^{-3}$ for SDSS (Blanton et al."941 2001. Yasuca et al.," 2001, Yasuda et al."942 2001)., 2001).943 The caveat is that no correction to total magnitudes is perfect. and the corrected: magnitudes will tend to have some surface brightness dependency as ds. demonstratec bv the Gaussian corrected Iuminositv function.," The caveat is that no correction to total magnitudes is perfect, and the corrected magnitudes will tend to have some surface brightness dependency as is demonstrated by the Gaussian corrected luminosity function."944" Even when isophotal magnitudes have been corrected: to. pseudo-tota magnitudes it is better to use a 1/l,,; or maximunt likelihood estimator which is a function of both AL and ji.", Even when isophotal magnitudes have been corrected to pseudo-total magnitudes it is better to use a $1/V_{max}$ or maximum likelihood estimator which is a function of both $M$ and $\mu$.945 While a good. understanding of visibility theory wil account for galaxies within the surface brightness limits. ealaxies with very low central surface brightness. but brieh total magnitudes can be missed.," While a good understanding of visibility theory will account for galaxies within the surface brightness limits, galaxies with very low central surface brightness, but bright total magnitudes can be missed."946 These are one source of mismatches between the estimates of a in dillerent surveys., These are one source of mismatches between the estimates of $\alpha$ in different surveys.947 Another reason for dillerent estimates is inhomogencities in the space density of galaxies., Another reason for different estimates is inhomogeneities in the space density of galaxies.948 Surveys looking at. dilferen parts of the sky will encounter variations in the Large Scale Structure., Surveys looking at different parts of the sky will encounter variations in the Large Scale Structure.949 Dwarf galaxies are seen over a smaller volume. so they can have large clustering corrections.," Dwarf galaxies are seen over a smaller volume, so they can have large clustering corrections."950" Dillerences in clustering corrections between different surveys will tend to bias a rather than 3"".", Differences in clustering corrections between different surveys will tend to bias $\alpha$ rather than $\phi^*$.951 We have presented a fitting function. for the Bivariate Brightness Distribution., We have presented a fitting function for the Bivariate Brightness Distribution.952 Vhis takes a similar form to a Schechter function in luminosity coupled: with a Ciaussian clistribution in surface brightness., This takes a similar form to a Schechter function in luminosity coupled with a Gaussian distribution in surface brightness.953 The Bivariate Drightness Funetion was fitted. to the recent. results from. Cross et al. (, The Bivariate Brightness Function was fitted to the recent results from Cross et al. (9542001) who constructed a BBD for a sample of 45.000 ealaxies [from the δαGIU.,"2001) who constructed a BBD for a sample of 45,000 galaxies from the 2dFGRS."955 The BBE fits the data well at the bright end. but poorly at the faint end.," The BBF fits the data well at the bright end, but poorly at the faint end."956 We compare the parameters of the fit to the de Jong Lacey (2000) results for late-tvpe spirals., We compare the parameters of the fit to the de Jong Lacey (2000) results for late-type spirals.957 While our results broadly agree there are dillerences which may provide clues toward understanding formation and. evolution. tracks οἱ dillerent galaxy tvpes., While our results broadly agree there are differences which may provide clues toward understanding formation and evolution tracks of different galaxy types.958" The BBE can be integrated to vield a total Iuminosity density of ji,=246E014105L. °"," The BBF can be integrated to yield a total luminosity density of $j_{b_j}=2.16\pm0.14\times10^8h\,L_{\odot}$ $^{-3}$."959" ""his agrees with the luminosity density calculated. from the data and demonstrates that unless the BBD shows sub-structure or a dramatic upturn bevonc the selection. boundaries the majority of the luminosity density in the local Universe has now been detected.", This agrees with the luminosity density calculated from the data and demonstrates that unless the BBD shows sub-structure or a dramatic upturn beyond the selection boundaries the majority of the luminosity density in the local Universe has now been detected.960 This paper has dealt. with the biases that occur if vou start with an isophotally selected sample and do not apply ux corrections. or apply a light loss correction without properly considering the volume correction.," This paper has dealt with the biases that occur if you start with an isophotally selected sample and do not apply any corrections, or apply a light loss correction without properly considering the volume correction."961 In both cases 1e Schechter parameters can be biased. and the luminosity ensity uncderestimated.," In both cases the Schechter parameters can be biased, and the luminosity density underestimated."962 Using a BBE we explore the impact of the limiting election isophote on classical measures of the Schechter function., Using a BBF we explore the impact of the limiting detection isophote on classical measures of the Schechter function.963" We demonstrate that if isophotal magnitudes are used then errors. of ~0.62 mags. z;NÓ~26% and Aa~0.04 are likelyAM at finn,=24.0 mag τν"," We demonstrate that if isophotal magnitudes are used then errors of $\Delta M^*_{b_j} \sim 0.62 $ mags, $\Delta \phi^* \sim 26\%$ and $\Delta \alpha \sim 0.04$ are likely at $\mu_{lim,b_j}=24.0$ mag $^{-2}$."964" if Gaussian. corrected magnitudes are used. these change to ~(38 mags. Ao~TA and Aa<0.01 are AM,likely at fine,=24.0 mag 7."," If Gaussian corrected magnitudes are used these change to $\Delta M^*_{b_j} \sim 0.38 $ mags, $\Delta \phi^* \sim 11\%$ and $\Delta \alpha < 0.01$ are likely at $\mu_{lim,b_j}=24.0$ mag $^{-2}$."965" JE total magnitudes can be recovered then the observed Luminosity. function will be correct. within the errors provided fni,724.0 mag 7."," If total magnitudes can be recovered then the observed luminosity function will be correct within the errors provided $\mu_{lim,b_j}>24.0$ mag $^{-2}$."966 Hence while the faint-end. slope. à. appears [zirlv robust to surface brightness issues both the A point and O° are highly. dependent.," Hence while the faint-end slope, $\alpha$, appears fairly robust to surface brightness issues both the $M^*$ point and $\phi^*$ are highly dependent."967 The range over which these parameters vary is fully consistent with the scatter in the published. values which come from a variety of surveys with cdillering selection criterion., The range over which these parameters vary is fully consistent with the scatter in the published values which come from a variety of surveys with differing selection criterion.968" These parameters produce & range in the Luminosity density. je, of 0.9.<jn,22.10HL. Alpe' again agreeing well with the range of published values (1.12.10hL.XMpc ! see Cross et al."," These parameters produce a range in the luminosity density, $j_{b_j}$, of $0.9 < j_{b_j} < 9692.2\times10^8h\,L_{\odot}$ $^{-3}$ again agreeing well with the range of published values $1.1-2.2\times10^8h\,L_{\odot}$ $^{-3}$ see Cross et al."970 2001)., 2001).971 When selection ellects are taken into account properly. the luminosity functions of recent surveys agree very well.," When selection effects are taken into account properly, the luminosity functions of recent surveys agree very well."972 The 20ECIUS and SDSS luminosity functions give the same AJ! and 6° values. but disagree on the values of a.," The 2dFGRS and SDSS luminosity functions give the same $M^*$ and $\phi^*$ values, but disagree on the values of $\alpha$."973 The differences in à are likely to be due to a combination of inpu catalogue incompleteness. filter used and dillerent clustering COLLECLLONS.," The differences in $\alpha$ are likely to be due to a combination of input catalogue incompleteness, filter used and different clustering corrections."974 85 suggests that if a deep. isophote is used. au total magnitudes are recovered. then traditional magnitucle dependent estimators can be usec.," 5 suggests that if a deep isophote is used and total magnitudes are recovered, then traditional magnitude dependent estimators can be used."975 However. future. work at higher redshilts or low redshift work on extremely fain ealaxies will include many galaxies close to the limits of the detection threshold. where surface brightness dependencies are strong.," However, future work at higher redshifts or low redshift work on extremely faint galaxies will include many galaxies close to the limits of the detection threshold, where surface brightness dependencies are strong."976 The BBD provides a framework that allows one to determine whether a bias is present. (see Fig. 3)), The BBD provides a framework that allows one to determine whether a bias is present (see Fig. \ref{fig:V_M}) )977 and to correct for it., and to correct for it.978" In addition. it produces parameters such as 3), and o, which could be used to place constraints on galaxy formation and evolution mocels."," In addition, it produces parameters such as $\beta_{\mu}$ and $\sigma_{\mu}$ which could be used to place constraints on galaxy formation and evolution models."979 Our conclusion is that after a quarter of a century we need to upgrade our representation of the space-cdensity of ealaxies to now include surface brightness., Our conclusion is that after a quarter of a century we need to upgrade our representation of the space-density of galaxies to now include surface brightness.980 The BBD and BBE provide an excellent. starting point ancl should. Lead to more reliable ancl Consistent measurements of the local luminosity density as well as providing new constraints on ealaxy formation mocels., The BBD and BBF provide an excellent starting point and should lead to more reliable and consistent measurements of the local luminosity density as well as providing new constraints on galaxy formation models.981This represents the change in the polar field from one cycle to the next.,This represents the change in the polar field from one cycle to the next.982 The correlation is expected to be stronger im this case. because the tilted BMRs first reverse the polar field of the old eycle and then built up that of the new cycle.," The correlation is expected to be stronger in this case, because the tilted BMRs first reverse the polar field of the old cycle and then built up that of the new cycle."983 The right panel of Figure 6 shows the relation between the maximum polar field at the end of a cycle and the strength of the next cycle., The right panel of Figure \ref{fig:corr_sn_pf} shows the relation between the maximum polar field at the end of a cycle and the strength of the next cycle.984 All correlations are significant at the p.<0.05 level., All correlations are significant at the $p\le 0.05$ level.985" From the results described in Section 3. we expect that for 7, the correlation between polar field and strength of the next cycle would be even stronger."," From the results described in Section 3, we expect that for $\eta_R=0$ the correlation between polar field and strength of the next cycle would be even stronger."986 Such a correlation is potentially relevant in connection with flux transport dynamo models (e.g.. 22).," Such a correlation is potentially relevant in connection with flux transport dynamo models \citep[e.g.,][]{Chatterjee04,Dikpati04}."987 Figure 7 shows the relationship between the total unsigned flux and the sunspot number., Figure \ref{fig:corr_sn_tf} shows the relationship between the total unsigned flux and the sunspot number.988 During activity maxima. many BMRs emerge and the polar fields are weak.," During activity maxima, many BMRs emerge and the polar fields are weak."989 Hence there is a strong correlation between the sunspot number and the total unsigned flux., Hence there is a strong correlation between the sunspot number and the total unsigned flux.990 During the minima fewer sunspots emerge. and the total surface flux has a stronger contribution from the polar region.," During the minima fewer sunspots emerge, and the total surface flux has a stronger contribution from the polar region."991 This leads to à weaker correlation between the total surface field and the sunspot numbers during solar minima., This leads to a weaker correlation between the total surface field and the sunspot numbers during solar minima.992 On longer timescales. the total flux is roughly proportional to the sunspot number.," On longer timescales, the total flux is roughly proportional to the sunspot number."993 Since the contribution to the field strength of the multipole of order / falls off with radius as 777. the lowest-order multipoles dominate the amplitude of the open heliospheric flux (?).. decreas," Since the contribution to the field strength of the multipole of order $l$ falls off with radius as $r^{-(l+2)}$, the lowest-order multipoles dominate the amplitude of the open heliospheric flux \citep{Wang02}."994eThe emergence of any given BMR may increase or the open flux. depending on whether its dipole moment vector is oriented so as to reinforce or reduce that of the pre-existing field (?)..," The emergence of any given BMR may increase or decrease the open flux, depending on whether its dipole moment vector is oriented so as to reinforce or reduce that of the pre-existing field \citep{Wang00}."995 The left panel of Figure 8 shows that the maximum value of open flux during a cycle. and the maximum sunspot number are strongly correlated., The left panel of Figure \ref{fig:corr_sn_of} shows that the maximum value of open flux during a cycle and the maximum sunspot number are strongly correlated.996 This results from the large number of BMRs emerging during the solar maximum period together with activity nesting. which generate a strong equatorial dipole field.," This results from the large number of BMRs emerging during the solar maximum period together with activity nesting, which generate a strong equatorial dipole field."997 On the other hand. the correlation is much weaker between the minimum value of open flux over a cycle and the minimum sunspot number.," On the other hand, the correlation is much weaker between the minimum value of open flux over a cycle and the minimum sunspot number."998 This is because the open flux during minimum phases ts dominated by the axial dipole field corresponding to the polar field that has accumulated over the cycle., This is because the open flux during minimum phases is dominated by the axial dipole field corresponding to the polar field that has accumulated over the cycle.999 Note that the maxima of the open flux usually occur 1—2 yrs after sunspot maximum (cf.?) while the minima of the open flux are almost in phase with the sunspot numbers., Note that the maxima of the open flux usually occur 1–2 yrs after sunspot maximum \citep[cf.][]{Wang02} while the minima of the open flux are almost in phase with the sunspot numbers.1000We assume that the cavity ancl Mercury. are almost spherical. this allows us to introduce the four small parameters ¢;: and also the parameter &=C/C. Le. the ratio between the polar inertial momentum of the core and of Mercury.,"We assume that the cavity and Mercury are almost spherical, this allows us to introduce the four small parameters $\epsilon_i$: and also the parameter $\delta=C_c/C$, i.e. the ratio between the polar inertial momentum of the core and of Mercury."1001 ει represents the polar flattening of Mercury. while e» is its equatorial ellipticity.," $\epsilon_1$ represents the polar flattening of Mercury, while $\epsilon_2$ is its equatorial ellipticity."1002 e; and e; have the same meaning for the cavity., $\epsilon_3$ and $\epsilon_4$ have the same meaning for the cavity.1003 If we assume the core of Mercury to be spherical. we should take €;=οι0. while €;=0 represents an axisvmmetric cavity.," If we assume the core of Mercury to be spherical, we should take $\epsilon_3=\epsilon_4=0$, while $\epsilon_4=0$ represents an axisymmetric cavity."1004 Henrard(2008). considered that the ellipsoid of inertia of the core and the mantle were aligned ancl proportional. the mathematical formulation was e;—ej and e;—e».," \citet{h08} considered that the ellipsoid of inertia of the core and the mantle were aligned and proportional, the mathematical formulation was $\epsilon_3=\epsilon_1$ and $\epsilon_4=\epsilon_2$."1005 Our parameters are gathered in Table 1.., Our parameters are gathered in Table \ref{tab:values}. .1006 We now introduce the two sets of Andover’s variables CXndover1926)... (0.gh.L.CH) and (GghobeGoHG). related respectively to the whole Mercury. and to its core.," We now introduce the two sets of Andoyer's variables \citep{a26}, $(l,g,h,L,G,H)$ and $(l_c,g_c,h_c,L_c,G_c,H_c)$, related respectively to the whole Mercury and to its core."1007 Phe angles (.A.9) are the Euler angles of the vector i. node of the equatorial plane over the plane perpendicular to the angular momentum A. the angles C/./) position the axis of least inertia fi with respect to 275.," The angles $(h,K,g)$ are the Euler angles of the vector $\vec{n_2}$ , node of the equatorial plane over the plane perpendicular to the angular momentum $\vec{N}$, the angles $(J,l)$ position the axis of least inertia $\vec{f_1}$ with respect to $\vec{n_2}$."1008 Correspondingly the angles (hi.dv.9g) are the Euler angles of the vector ni. node of the equatorial plane over the plane perpendicular to the angular momentum of the pseudo-core IN. and. (ονο). position the axis of least inertia with respect to n5.," Correspondingly the angles $(h_c,K_c,g_c)$ are the Euler angles of the vector $\vec{n^c_2}$, node of the equatorial plane over the plane perpendicular to the angular momentum of the pseudo-core $\vec{N_c}$, and $(J_c,l_c)$ position the axis of least inertia with respect to $\vec{n^c_2}$ ."1009 Figure 2. shows a schematic view of all the reference frames and relevant angles., Figure \ref{fig:bigfig} shows a schematic view of all the reference frames and relevant angles.1010" Phe variables are (hog!) and (higf) and the corresponding momenta (ff=NcosNK. €=IN. L=Ncos J) and (H1,=Nocos νε. Cy=NT. L,=N'cosJ, )."," The variables are $(h,g,l)$ and $(h_c,g_c,l_c)$ and the corresponding momenta $H=N\cos K$, $G=N$, $L=N\cos J$ ) and $H_c=N^c \cos K_c$ , $G_c=N^c$, $L_c=N^c \cos J_c$ )."1011 Eixpressed in Anclover’s variables the components of N and INC are: We can now straightforwardly derive the Hamiltonian Hy of the free rotation of Mercury. using Andover's variables and changing the sign of Nc to take the minus signof the PoincaréHough equationsinto account (I2.20)).," Expressed in Andoyer's variables the components of $\vec{N}$ and $\vec{N^c}$ are: We can now straightforwardly derive the Hamiltonian $\mathcal{H}_1$ of the free rotation of Mercury, using Andoyer's variables and changing the sign of $\vec{N^c}$ to take the minus signof the Poincaré–Hough equationsinto account \ref{equ:ph2}) )."1012 We also linearize the Hamiltonian with respect to thesmall parameters e; (their orders ofmagnitude being about 10. 77). and get:," We also linearize the Hamiltonian with respect to thesmall parameters $\epsilon_i$ (their orders ofmagnitude being about $10^{-5}$ ), and get:"1013Ihela.,",."1014. This average drilt vanishes in the equatorial plane. ancl has opposite sigus$ in opposite (rotational) henispheres.," This average drift vanishes in the equatorial plane, and has opposite signs in opposite (rotational) hemispheres."1015 The αν given by equation (32)) is superimposed on the assumed rigid corotation., The drift given by equation \ref{Deltav1}) ) is superimposed on the assumed rigid corotation.1016 The dift does not correspond to a modification of the angular velocity., The drift does not correspond to a modification of the angular velocity.1017 To see this consider a perturbation. dw. in the angular speed of (he magnetosphere.," To see this consider a perturbation, $\delta\omega$, in the angular speed of the magnetosphere."1018 This would give a perturbation δω rsin&inaziunmuthal velocity: (his has a maximum in the equatorial plane. has the same sien in both hemispheres aud is independent of rotational phase.," This would give a perturbation $\delta \omega\, r\sin\theta$ in azimuthal velocity; this has a maximum in the equatorial plane, has the same sign in both hemispheres and is independent of rotational phase."1019 The velocity given by equation (32)) satisfies none of these conditions., The velocity given by equation \ref{Deltav1}) ) satisfies none of these conditions.1020 The magnetosphere cannot be in rigid rotation al any angularvelocity., The magnetosphere cannot be in rigid rotation at any angularvelocity.1021" Figures 2.. 3. and 4 show plots of Av,,. Ary and Av,. respectively. in units of wr as [functions ofwl with /=0 at o=0."," Figures \ref{fig-SC28v_ph}, , \ref{fig-SC28v_th} and \ref{fig-SC28v_r} show plots of $\Delta v_\phi$, $\Delta v_\theta$ and $\Delta v_r$ , respectively, in units of $\omega r$ as functions of$\omega t$ with $t = 0$ at $\phi = 0$."1022 Each curve represents values for 6=Q. 107. 20° or 30° for a=30°.," Each curve represents values for $\theta = 0^\circ$, $10^\circ$, $20^\circ$ or $30^\circ$ for $\alpha = 30^\circ$."1023 The correction in square brackets in equation (32)) is included. so that the temporal variations include a term varying sinusoidallv as ó—wl wilh a correction tern that includes a variation as 2(6—«/).," The correction in square brackets in equation \ref{Deltav1}) ) is included, so that the temporal variations include a term varying sinusoidally as $\phi - \omega t$ with a correction term that includes a variation as $2(\phi - \omega t)$."1024" The components Ar,, and Av, are odd functions of cos. implvingthat thev have opposite signs in opposite hemispheres. whereas Ary is an even function of eos8."," The components $\Delta v_\phi$ and $\Delta v_r$ are odd functions of ${\rm cos}\, \theta$, implyingthat they have opposite signs in opposite hemispheres, whereas $\Delta v_\theta$ is an even function of ${\rm cos}\, \theta$."1025 The pulse window corresponds to a small range of © aud 9. aid. because an observer can see the drift only within (his window. only a correspondingly small range in figures 2.. 3 and 4 is relevant to observations.," The pulse window corresponds to a small range of $\phi$ and $\theta$, and because an observer can see the drift only within this window, only a correspondingly small range in figures \ref{fig-SC28v_ph}, \ref{fig-SC28v_th} and \ref{fig-SC28v_r} is relevant to observations."1026" This range is model-dependent. ancl in most models 1 corresponds to a range «Nó about ©=0° and a range AG, of order Ao. about a line of sight. 8. close to a."," This range is model-dependent, and in most models it corresponds to a range $\Delta\phi$ about $\phi = 0^\circ$ and a range $\Delta\theta$, of order $\Delta\phi$ , about a line of sight, $\theta$, close to $\alpha$."1027 The relevant regions in figures 2.. 2. and 4 are near &/=0 for the solid curves.," The relevant regions in figures \ref{fig-SC28v_ph}, \ref{fig-SC28v_th} and \ref{fig-SC28v_r} are near $\omega t = 0$ for the solid curves."1028 The three curves for 9=a30° are replotted in ligure 5.., The three curves for $\theta = \alpha = 30^\circ$ are replotted in figure \ref{fig-AsymmetryAlign}.1029 As the variation with 8 is relatively small. (he observable (relative to corotation) drifts correspondto the range Ao about wf=0in figure 5..," As the variation with $\theta$ is relatively small, the observable (relative to corotation) drifts correspondto the range $\Delta\phi$ about $\omega t = 0$in figure \ref{fig-AsymmetryAlign}. ."1030" The components Ary and Ar, ave zero at οἱ= 0. implving that the observable drift is in the ©direction."," The components $\Delta v_\theta$ and $\Delta v_r$ are zero at $\omega t = 0$ , implying that the observable drift is in the $\phi$direction."1031 The magnitude of this drift may be estimated by selling 8— a. o—ut0 in equation (32)).," The magnitude of this drift may be estimated by setting $\theta = \alpha$ , $\phi - \omega t = 0$ in equation \ref{Deltav1}) )."1032 This implies an observable ciii, This implies an observable drift1033obtained the JITIS magnitudes of the stars aud companion candidates from the SUARP-T images (because the objects are best separated on those nuage. see Fig.,"obtained the JHK magnitudes of the stars and companion candidates from the SHARP-I images (because the objects are best separated on those image, see Fig."1034 d 2) by normal aperture photometry with an aperture size just sanall chough to exclude the companion. compared to the flux standard IIR 8278(van der Blick et al.," 1 2) by normal aperture photometry with an aperture size just small enough to exclude the companion, compared to the flux standard HR 8278 (van der Bliek et al."1035 1996) observed in JUS iuunediatelv after the respective scieuce target., 1996) observed in JHK immediately after the respective science target.1036 Thei the maguitudes for the companion candidates are obtained roni the backeround-subtracted flux ratio with their respective primary.," Then, the magnitudes for the companion candidates are obtained from the background-subtracted flux ratio with their respective primary."1037 The resulting photometric data are isted in Table 2., The resulting photometric data are listed in Table 2.1038 We display the SILTARP-I JITIS images of ΠΟ 199113 aud its close (17) companion candidate in Fig., We display the SHARP-I JHK images of HD 199143 and its close $1^{\prime \prime}$ ) companion candidate in Fig.1039 1 together with the Dec. 2000 Soff image. where the companion caudidate (called ΠΟ 199113 Bin JBOL) was iargimallv detected.," 1 together with the Dec. 2000 SofI image, where the companion candidate (called HD 199143 B in JB01) was marginally detected."1040 Our data on ΠΟ 199113 A aeree well with those iu JBOL. within our more conservative precision of 0.1 mae (as we observed less standards). while ΠΟ 199113 D is 0.3 mae brighter m our data than iun JDOl. possibly due to problems in JBOL with the subtraction of the xieht background due to the PSF of ΠΟ 199118 A(R. Javawardhana. priv.," Our data on HD 199143 A agree well with those in JB01, within our more conservative precision of $\pm 0.1$ mag (as we observed less standards), while HD 199143 B is 0.3 mag brighter in our data than in JB01, possibly due to problems in JB01 with the subtraction of the bright background due to the PSF of HD 199143 A (R. Jayawardhana, priv."1041 coun.}., comm.).1042 The IR colors Jff=O0.2401danudi W=0.0+0.1 nag (Table 2) for ΠΟ 199113 A are cousisteut with its spectral type Fs aud its optical colors. for negligible oreground extinction.," The IR colors $J-H=0.2 \pm 0.1$ and $H-K=0.0 \pm 0.1$ mag (Table 2) for HD 199143 A are consistent with its spectral type F8 and its optical colors, for negligible foreground extinction."1043 The companion candidate to TD 199113 is redder than the primary. either because it ds COSS lnassive and/or more reddened. muuelv bv either cirectuustellar material and/or other foreground extinction.," The companion candidate to HD 199143 is redder than the primary, either because it is less massive and/or more reddened, namely by either circumstellar material and/or other foreground extinction."1044 Iu case of nceligible absorption. as towards the primary. he JIIIN colors of the companion candidate would be consistent within the errors with a nud-I to mid-AL type dwarf.," In case of negligible absorption, as towards the primary, the JHK colors of the companion candidate would be consistent within the errors with a mid-K to mid-M type dwarf."1045 If the companion caudidate is at the same distance as the primary and also on the zero-age main-sequence. than its absolute JTS magnitudes are best consisteut with an carly ALtype chart.," If the companion candidate is at the same distance as the primary and also on the zero-age main-sequence, than its absolute JHK magnitudes are best consistent with an early M-type dwarf."1046 Heuce. our spectral type classification is MO-2," Hence, our spectral type classification is M0-2."1047 IID DPS9SG:& and its close (27) companion caudidate (called IID Dand5Nο D iu JBOL) are shown in Fig., HD 358623 and its close $2^{\prime \prime}$ ) companion candidate (called HD 358623 B in JB01) are shown in Fig.1048 2., 2.1049 The primary star the companion candidate are also wel separated and detected iu our Sofl images. so that we cau obtain photometric and astrometric data ποια both the Sofl and the SITARP-I images.," The primary star and the companion candidate are also well separated and detected in our SofI images, so that we can obtain photometric and astrometric data from both the SofI and the SHARP-I images."1050 In case of the Sofl images. we used the faint IST standard stars S-361-D and S-751-C (Persson et al.," In case of the SofI images, we used the faint HST standard stars S-361-D and S-754-C (Persson et al."1051" 1998). and the data aeree well with our SITARP-I IIland magnitudes,"," 1998), and the data agree well with our SHARP-I H-band magnitudes."1052 Our data ou ΠΟ 358623 A and D also agree well with those iu JBOL., Our data on HD 358623 A and B also agree well with those in JB01.1053 ITeuce. we lave 10 Indication for variability.," Hence, we have no indication for variability."1054" The IR MN jII—0.7zO0.laud/I W=O140.1 nae Hx IID3-- A areconsistent with its spectral ype τον. ""nmefor de extinction."," The IR colors $J-H=0.7 \pm 0.1$ and $H-K=0.1 \pm 0.1$ mag for HD 358623 A are consistent with its spectral type K7-M0, for negligible extinction."1055 There is no reddening in the near IR in hese two stars. consistent with he spectral energv distribution shown im vdAOL. where some moderate IR excess is seen ouly in the thermal IR.," There is no reddening observed in the near IR in these two stars, consistent with the spectral energy distribution shown in vdA01, where some moderate IR excess is seen only in the thermal IR."1056 The IR colors of the companion candidate TD 358623 D are very similar as for the primary IID 358623. but also consistent with any spectral type between nid-I& aud iid-AL," The IR colors of the companion candidate HD 358623 B are very similar as for the primary HD 358623, but also consistent with any spectral type between mid-K and mid-M."1057 Furthermore. the companion candidate is oulv slightly züuter (less than 2 mae) thanthe primary. so that it could ο a shehtly redder stellar secoudirv with slightly less uass and slightly later spectral type (early to τςΑΓ).," Furthermore, the companion candidate is only slightly fainter (less than 2 mag) than the primary, so that it could be a slightly redder stellar secondary with slightly less mass and slightly later spectral type (early to mid-M)."1058" A spectral type MO-3is best consistent with its absolute ανας, iaguitudes. when assuming that is has the sale distance as TD 199113 A and that it is also located ou the zero-age main-sequence."," A spectral type M0-3 is best consistent with its absolute JHK-band magnitudes, when assuming that is has the same distance as HD 199143 A and that it is also located on the zero-age main-sequence."1059 This classification will be refined by au ITE-baud spectrum (Sect., This classification will be refined by an H-band spectrum (Sect.1060 5)., 5).1061 Whether the two conrpanon candidates are truclv bound companions or (possibly reddened) background stars can be found out best by spectroscopy aud proper motion., Whether the two companion candidates are truely bound companions or (possibly reddened) background stars can be found out best by spectroscopy and proper motion.1062 Ata period of 1.57 bh RAJ1L089 is clearly below the orbital period σαρ of CVs.,At a period of 1.57 h RXJ1039 is clearly below the `orbital period gap' of CVs.1063 The very large amplitude of the principal photometric modulation is the signature of a reflection cllect., The very large amplitude of the principal photometric modulation is the signature of a reflection effect.1064 There are no detectable eclipses. which implies an inclination less than about 65°. but the inclination cannot be much less than this for. with its peak-to-peak range of 1.1 mag. 11NJ1039 has among the largest amplitude modulations observed for CVs and detached white ΑνναΣΙ dwacrl binaries (scc. e.g... Chen et al.," There are no detectable eclipses, which implies an inclination less than about $^\circ$, but the inclination cannot be much less than this for, with its peak-to-peak range of 1.1 mag, RXJ1039 has among the largest amplitude modulations observed for CVs and detached white dwarf/M dwarf binaries (see, e.g., Chen et al."1065 1995)., 1995).1066 Phe implication is that the primary in 1.)1099 is very hot. and this in turn suggests that it underwent a nova eruption within the past decade or so.," The implication is that the primary in RXJ1039 is very hot, and this in turn suggests that it underwent a nova eruption within the past decade or so."1067 Phe known novae with periods below the period gap (RW UAL. GO Mus. CP Pup. VI974 (νο RS Car) include representatives from ast and slow classes of light curves. with eruption ranges around 10 - 13 mag.," The known novae with periods below the period gap (RW UMi, GQ Mus, CP Pup, V1974 Cyg, RS Car) include representatives from fast and slow classes of light curves, with eruption ranges around 10 - 13 mag."1068 With a quicscent magnitude of 18.0. UNJ1039 would probably therefore not have reached: naked eve brightness and would be most likely to have been found or missed in photographic surveys (which in any case have »ven reduced in completeness in the past two decades).," With a quiescent magnitude of 18.0, RXJ1039 would probably therefore not have reached naked eye brightness and would be most likely to have been found – or missed – in photographic surveys (which in any case have been reduced in completeness in the past two decades)."1069 At (=2537. b=45. RAJLO89 is [ar from the direction of the Galactic Centre. towards which most photographie surveys or novae have been concentrated.," At $\ell = 253^\circ$, $b = 45^\circ$, RXJ1039 is far from the direction of the Galactic Centre, towards which most photographic surveys for novae have been concentrated."1070 RAJLO39 is therefore a ime cancdidate for first membership of a class of relatively recent ‘overlooked’ novae., RXJ1039 is therefore a prime candidate for first membership of a class of relatively recent `overlooked' novae.1071 The ROSAT catalogue of X-tav sources may contain other such wails: their eventual identification will give an indication of the completeness of optical searches in the past. decades., The ROSAT catalogue of X-Ray sources may contain other such waifs; their eventual identification will give an indication of the completeness of optical searches in the past decades.1072 We have searched the available data archives to see if any observations in the vicinity of 11X1039 were made in the past 10 vears., We have searched the available data archives to see if any observations in the vicinity of RXJ1039 were made in the past 10 years.1073 No observations were found. although we note the proximity (775 aresee) of a small galaxy (231ASXI J1039517-050640).," No observations were found, although we note the proximity $\sim$ 75 arcsec) of a small galaxy (2MASXi J1039517-050649)."1074 The ET (of intensities) of our entire data set (with the mean of each run subtracted). prewhitened at the orbital period and its first harmonic. is shown in Figure 2..," The FT (of intensities) of our entire data set (with the mean of each run subtracted), prewhitened at the orbital period and its first harmonic, is shown in Figure \ref{ftrxj1039}."1075 We have not displayed the lowest frequencies. where there is à broad. low amplitude feature centred on the orbital frequeney. which we attribute to sidebands that occur because of small non-periodic modulation of the signal.," We have not displayed the lowest frequencies, where there is a broad low amplitude feature centred on the orbital frequency which we attribute to sidebands that occur because of small non-periodic modulation of the signal."1076 Various combinations of runs show two persistent prominent. peaks in the vicinity of 1933 s and 723 s. In the ET of the complete data. set these appear at 1932.5 s with an zunplitude of. and at 721.87 s with an amplitude of144.," Various combinations of runs show two persistent prominent peaks in the vicinity of 1933 s and 723 s. In the FT of the complete data set these appear at 1932.5 s with an amplitude of, and at 721.87 s with an amplitude of."1077. Phere are also aliases to these periods κος Table 2 and. Figure 3..., There are also aliases to these periods – see Table \ref{tab2} and Figure \ref{ftburxj1039}.1078 Expressed as frequencies. the two prominent modulations are at fi = 517.5 plz and fo = 1383.0 11 respectively. ancl we notice that fo 2 fj = 348.0 pelle z20. where © is the orbital frequeney. (176.44. plz).," Expressed as frequencies, the two prominent modulations are at $f_1$ = 517.5 $\mu$ Hz and $f_2$ = 1383.0 $\mu$ Hz respectively, and we notice that $f_2$ – 2 $f_1$ = 348.0 $\mu$ Hz $\approx 2 \Omega$, where $\Omega$ is the orbital frequency (176.44 $\mu$ Hz)."1079 From comparison with the optical modulations seen in intermediate polars (LPs - e.g. Patterson (1994)) this suggests the identification of fj with the frequency wO and. fo with δω. where w is the spin period. of the white dwarf primary.," From comparison with the optical modulations seen in intermediate polars (IPs - e.g. Patterson (1994)) this suggests the identification of $f_1$ with the frequency $\omega - \Omega$ and $f_2$ with $\omega$, where $\omega$ is the spin period of the white dwarf primary."1080 Εις receives support from the existence in the first (wo of our runs of a signal near w which shows in the total PT as à modulation at 1243.7 s (ie. 692.7 pillz)., This receives support from the existence in the first two of our runs of a signal near $\omega$ which shows in the total FT as a modulation at 1443.7 s (i.e. 692.7 $\mu$ Hz).1081 We also find à modulation in the ET at 640.3 8 (1562 pllz) that we identify as à 2o | O modulation., We also find a modulation in the FT at 640.3 s (1562 $\mu$ Hz) that we identify as a $\omega$ + $\Omega$ modulation.1082 There is à peak of power in the FP near ac | O modulation (sce Figure 2)) but we cannot find a convincing amplitude spike that coincides with it or its alias: there is no evidence in the FEE fora δω © modulation., There is a peak of power in the FT near a $\omega$ + $\Omega$ modulation (see Figure \ref{ftrxj1039}) ) but we cannot find a convincing amplitude spike that coincides with it or its alias; there is no evidence in the FT for a $\omega$ – $\Omega$ modulation.1083 Only the set of frequencies that we have discussed. here produces a sensible mocel choice of any of the aliases of these selected modulations does not lead to any numerical relationships that match what is recognised as typical LP behaviour., Only the set of frequencies that we have discussed here produces a sensible model – choice of any of the aliases of these selected modulations does not lead to any numerical relationships that match what is recognised as typical IP behaviour.1084 We have not attached: uncertainties to the frequency measurements: many of the mocdulations do not rise much above the noise (see Figure 3)) and will be inlluenced. in an unknown and non-rancom wav by the noise., We have not attached uncertainties to the frequency measurements; many of the modulations do not rise much above the noise (see Figure \ref{ftburxj1039}) ) and will be influenced in an unknown and non-random way by the noise.1085 For our strongest (jw QO) modulation we find from an ET of runs A- Band € 517.4 pillz (2.5343) and from runs €. D and I 516.7 plz D.," For our strongest $\omega - \Omega$ ) modulation we find from an FT of runs A, B and C 517.4 $\mu$ Hz ) and from runs C, D and E 516.7 $\mu$ Hz )."1086 In the weaker signals errors of frequency of 1 plz or greater can therefore be expected., In the weaker signals errors of frequency of 1 $\mu$ Hz or greater can therefore be expected.1087 This suite of modulations gives strong evidence that ΗΝ.1020 is an LP with a spin period of 1444 s. accreting on to two magnetic poles. where the strongest optical signal is," This suite of modulations gives strong evidence that RXJ1039 is an IP with a spin period of 1444 s, accreting on to two magnetic poles, where the strongest optical signal is"1088lU we set CHO)=UG). and further assume that the subclumps are much smaller than the parent halo. then the power spectrum reduces to the sum of the power spectra of the individual components: This shows that whether or not the substructure component dominates the small scale. power depends. on the fraction of mass in the subelumps. and on how much more dense the subclumps are compared. to the smooth component.,"If we set $U_c(k)=U(k)$, and further assume that the subclumps are much smaller than the parent halo, then the power spectrum reduces to the sum of the power spectra of the individual components: This shows that whether or not the substructure component dominates the small scale power depends on the fraction of mass in the subclumps, and on how much more dense the subclumps are compared to the smooth component."1089 To see this a little more clearly. consider a specific example.," To see this a little more clearly, consider a specific example."1090 Suppose that the density runs are Ciaussians. ancl that all the IN subelumps have the same mass m; and size D.," Suppose that the density runs are Gaussians, and that all the $N$ subclumps have the same mass $m_i$ and size $R_i$."1091 WR denotes the characteristic scale of the density runs of the smooth component. then (=οχρίLRT/2) and ny=expl E7472).," If $R$ denotes the characteristic scale of the density runs of the smooth component, then $U=\exp(-k^2R^2/2)$ and $u_i = \exp(-k^2R_i^2/2)$ ."1092 Ehe subclumps dominate the power if Nm;pyra(uMpyUG.," The subclumps dominate the power if $N(m_i/\bar\rho)^2 u_i^2 > (M/\bar\rho)^2\,U(k)^2$."1093" ""Phis happens on scales where h2InCM1)|InN]""a1Ry]. where we have set AM,=Nm;."," This happens on scales where $k^2 > [2\ln (M/M_{cl})+\ln N]/R^2/[1 - (R_i/R)^2]$, where we have set $M_{cl} \equiv Nm_i$."1094 Note that if the total mass in subclumps is small. AZ4/AM l. then the subelumps dominate the power onlv at very large A: if this mass is divided up among many subclumps. the subclumps only dominate at. even higher A.," Note that if the total mass in subclumps is small, $M_{cl}/M\ll 1$ , then the subclumps dominate the power only at very large $k$; if this mass is divided up among many subclumps, the subclumps only dominate at even higher $k$."1095 For a fixed mass ratio.the exact scale on which the subclumps dominate depends on the size ratio £7/4.," For a fixed mass ratio,the exact scale on which the subclumps dominate depends on the size ratio $R_i/R$."1096 This suggests that a feature in P(&) at large & may provide information about the nature of the subclumps., This suggests that a feature in $P(k)$ at large $k$ may provide information about the nature of the subclumps.1097 Simulations suggest that ΛΙ~10 and ΕνXLF/10. which vields a critical value of &ziνο|InNA.," Simulations suggest that $M/M_{cl}\sim 10$ and $R_i\lsim R/10$, which yields a critical value of $k \gsim \sqrt{6 + \ln N}/R$."1098 BH is interesting that this is just bevond the reach of current simulations., It is interesting that this is just beyond the reach of current simulations.1099 Fig., Fig.1100 1l illustrates this with slightlv more realistic varro. Frenk White (1997) density profiles.," \ref{pksubclumps} illustrates this with slightly more realistic Navarro, Frenk White (1997) density profiles."1101 The various curves show AA)=4x4PCA) for models in which he density field is mace up of Poisson distributed NEW alos.," The various curves show $\Delta(k)\equiv 4\pi k^3\,P(k)$ for models in which the density field is made up of Poisson distributed NFW halos."1102 “The parent halos are truncated at their virial radii ray defined so that the average density within ry is 200 imes that ofthe background., The parent halos are truncated at their virial radii $r_{\rm vir}$ defined so that the average density within $r_{\rm vir}$ is 200 times that of the background.1103 We set the NEW core radius o be anpy=Olea.," We set the NFW core radius to be $a_{\rm NFW} =0.1\,r_{\rm vir}$."1104 The dashed line (same in both xuels) shows A(A) i£ there is no substructure., The dashed line (same in both panels) shows $\Delta(k)$ if there is no substructure.1105 We then assumed that a fraction f of the mass of each parent halo was in subclumps. each of mass m.," We then assumed that a fraction $f$ of the mass of each parent halo was in subclumps, each of mass $m$."1106 Lor simplicity. we assumed that. the distribution of subelumps. around the halo center was given by the same NEW form. and. that the distribution of particles around cach subelump center was also NEW. with core radius expw=0.157...," For simplicity, we assumed that the distribution of subclumps around the halo center was given by the same NFW form, and that the distribution of particles around each subclump center was also NFW, with core radius $a_{\rm NFW}=0.1\,r_{\rm sc}$."1107 Phe value of ri; was set by requiring that the average density within the subelumps equal 94; times the background density., The value of $r_{\rm sc}$ was set by requiring that the average density within the subclumps equal $\delta_{\rm sc}$ times the background density.1108 The solid. curves in the panel on the left show equation. (21)) when f= Q.l. m=Q.01A/7 and ὃν=500 (lower) and à.=5000 (upper).," The solid curves in the panel on the left show equation \ref{pk}) ) when $f=0.1$ , $m=0.01M$ and $\delta_{\rm sc}=500$ (lower) and $\delta_{\rm sc}=5000$ (upper)."1109 This shows that. all other things being equal. denser subclumps contribute more power.," This shows that, all other things being equal, denser subclumps contribute more power."1110 The solid curves in the panel on the right show results when f=0.9. δν=500 and m/AL=0.1 and m/AJ=0.01 (departure from dashed. curve apparent at high. and still higher. 4).," The solid curves in the panel on the right show results when $f=0.9$, $\delta_{\rm sc}=500$ and $m/M=0.1$ and $m/M=0.01$ (departure from dashed curve apparent at high and still higher $k$ )."1111 Comparison with the panel on the left shows that increasing the fraction of mass in subclumps. f. increases the small scale power.," Comparison with the panel on the left shows that increasing the fraction of mass in subclumps, $f$, increases the small scale power."1112 The scale on which this increase becomes apparent depends on (mnMy)(Guluu)yt5 the typical radii of the subelumps.," The scale on which this increase becomes apparent depends on $(m/M)^{1/3}(\delta_{\rm vir}/\delta_{\rm sc})^{1/3}$, the typical radii of the subclumps."1113 lligher order correlations at à given small scale are dominated more strongly by the one-halo term than is the two point function. so the οσοι of substructure on the one-halo term is of great interest for higher order correlations.," Higher order correlations at a given small scale are dominated more strongly by the one-halo term than is the two point function, so the effect of substructure on the one-halo term is of great interest for higher order correlations."1114 ln this section we will consider the elfect of substructure on the 3-point correlation function., In this section we will consider the effect of substructure on the 3-point correlation function.1115 The extension to higher orders is obvious., The extension to higher orders is obvious.1116 We begin with the expansion of the density p(r) given in equation (2))., We begin with the expansion of the density $\rho(r)$ given in equation \ref{rhoexpansion}) ).1117 The 3-point function in real space is then defined by ]t is convenient to consider the expressions for higher order correlations in Fourier space. since we will find that in the limiting case of interest the only additional term due to substructure is a term involving products of the substructure profile.," The 3-point function in real space is then defined by It is convenient to consider the expressions for higher order correlations in Fourier space, since we will find that in the limiting case of interest the only additional term due to substructure is a term involving products of the substructure profile."1118 The Fourier transform. of the 3-point. function is the bispectrum £2 delined by where the Deltafunction indicates that &j|ke , The Fourier transform of the 3-point function is the bispectrum $B$ defined by where the Deltafunction indicates that ${\bm k_1 + \bm k_2 + \bm k_3}=0$.1119Lt is the sum of contributions from triplets which are all in the smooth component. triplets all in the same halo. pluscontributions from the various cross terms.," It is the sum of contributions from triplets which are all in the smooth component, triplets all in the same halo, pluscontributions from the various cross terms."1120 Writing all the different terms explicitly gives, Writing all the different terms explicitly gives1121from which we measure the physical properties of the component stars to high precision.,from which we measure the physical properties of the component stars to high precision.1122 This study highlights the many possibilities opening up for the study of eclipsing binaries of all types from the increasing number of long-term variability surveys. in particular those which have the photometric precision to search for extrasolar planetary transits.," This study highlights the many possibilities opening up for the study of eclipsing binaries of all types from the increasing number of long-term variability surveys, in particular those which have the photometric precision to search for extrasolar planetary transits."1123 The object 02460 was discovered to be an eclipsing binary system by Strohmeier&Knigge(1961). who found that an orbital period of dd satistied their photographie observations but yielded a light curve with no detectable secondary minimum.," The object 460 was discovered to be an eclipsing binary system by \citet{StrohmeierKnigge61}, who found that an orbital period of d satisfied their photographic observations but yielded a light curve with no detectable secondary minimum."1124 The first photoelectric study of CCet (Morrison&Morrison1968) revised the orbital period to dd. resulting in a light curve containing rather similar primary and secondary minima.," The first photoelectric study of Cet \citep{MorrisonMorrison68aj} revised the orbital period to d, resulting in a light curve containing rather similar primary and secondary minima."1125 Morrison Morrison analysed their light curve using the method of RussellMerrill (1952). finding that the two stars were very similar.," Morrison Morrison analysed their light curve using the method of \citet{RussellMerrill52book}, finding that the two stars were very similar."1126 Popper(1969). announced that both components of CCet had a metallic-lined nature., \citet{Popper69baas} announced that both components of Cet had a metallic-lined nature.1127 This was independently discovered by Young (Popper1971)., This was independently discovered by Young \citep{Popper71apj}.1128. The first — and so far only — time-resolved spectroscopic study was presented by Popper(1971). based on photographic material., The first – and so far only – time-resolved spectroscopic study was presented by \citet{Popper71apj} based on photographic material.1129 Popper combined his spectroscopic results with the photometric elements of Morrison&Morrison(1968) to make the first determination of the masses and radii of the component stars of CCet. finding the two objects to have similar properties (masses 1.75 and aand radii 1.95 and 9).," Popper combined his spectroscopic results with the photometric elements of \citet{MorrisonMorrison68aj} to make the first determination of the masses and radii of the component stars of Cet, finding the two objects to have similar properties (masses 1.75 and and radii 1.95 and )."1130 Popper gave a rough spectral classification of A2m + FOm. However. an analysis of new photometric data for CCet led Srivastava&Padalia(1975) to radii which were quite different (2.and 9.," Popper gave a rough spectral classification of A2m + F0m. However, an analysis of new photometric data for Cet led \citet{SrivastavaPadalia75apss} to radii which were quite different (2.14 and )."1131 Okazaki(1978) revisited the photometric data obtainedI4 by both Morrison&(1968) and Srivastava&Padalia(1975) and resolved this discrepancy in favour of Popper's results (radii. 1.96 and 12.3). in the process finding that CCet has a small but significant orbital eccentricity.," \citet{Okazaki89apss} revisited the photometric data obtained by both \citet{MorrisonMorrison68aj} and \citet{SrivastavaPadalia75apss} and resolved this discrepancy in favour of Popper's results (radii 1.96 and ), in the process finding that Cet has a small but significant orbital eccentricity."1132 However. a substantial disparity between the radii of the two components was resurrected by Ramellaetal.(1980). through an analysis of the data obtained by Srivastava&Padalia(1975)..," However, a substantial disparity between the radii of the two components was resurrected by \citet{Ramella+80apss}, through an analysis of the data obtained by \citet{SrivastavaPadalia75apss}."1133" These authors used. for the first time in the case of CCet. ""modern light curve modelling techniques” to find radii of 2.10 and1."," These authors used, for the first time in the case of Cet, `modern light curve modelling techniques' to find radii of 2.10 and."1134741%... These values were supported by Srivastava(1987). who reanalysed the data from Srivastava&Padalia(1975) to find values of 2.13 and.," These values were supported by \citet{Srivastava75apss}, who reanalysed the data from \citet{SrivastavaPadalia75apss} to find values of 2.13 and."1135. Since the apparent resolution of varied measurements of the radii of the component stars of CCet in favour of a notable difference between the two objects. there has been only a little interest in this binary system.," Since the apparent resolution of varied measurements of the radii of the component stars of Cet in favour of a notable difference between the two objects, there has been only a little interest in this binary system."1136 Srivastava(1988) presented a period study of CCet. finding tive changes in period but providing no clear evidence for their existence.," \citet{Srivastava87apss} presented a period study of Cet, finding five changes in period but providing no clear evidence for their existence."1137 Finally. one spectrum of the object was obtained by Glazunovaetal.C008)... who measured the rotational velocities of both components.," Finally, one spectrum of the object was obtained by \citet{Glazunova+08aj}, who measured the rotational velocities of both components."1138 Their observation was taken at an orbital phase of 0.957. where the spectral lines of the two components would be partially overlapping. so it is no surprisingthat their results are inconsistent with the values we finc in refseczatm..," Their observation was taken at an orbital phase of 0.957, where the spectral lines of the two components would be partially overlapping, so it is not surprisingthat their results are inconsistent with the values we find in \\ref{sec:atm}."1139 The basic observational properties of CCet are given in reftab:basedata.., The basic observational properties of Cet are given in \\ref{tab:basedata}.1140 Throughout this work we refer to the primary star as AA and the secondary star as BB. AA is hotter. larger and more massive than BB. Spectroscopic observations were obtained in 2002 October using the mm Isaae. Newton Telescope (INT) on La Palma.," Throughout this work we refer to the primary star as A and the secondary star as B. A is hotter, larger and more massive than B. Spectroscopic observations were obtained in 2002 October using the m Isaac Newton Telescope (INT) on La Palma."1141 The 500mmm camera of the Intermediate Dispersion Spectrograph (DS) was equipped with a holographic + grating., The mm camera of the Intermediate Dispersion Spectrograph (IDS) was equipped with a holographic $^{-1}$ grating.1142 An EEV 22k CCD was used and exposure times were ss. From measurements of the full width at half maximum (FWHM) of are lines taken for wavelength calibration we estimate that the resolution is (1.8 pixels)., An EEV $\times$ 2k CCD was used and exposure times were s. From measurements of the full width at half maximum (FWHM) of arc lines taken for wavelength calibration we estimate that the resolution is (1.8 pixels).1143 A total of 38 spectra were taken covering the intervalAA.. with estimated signal to noise ratios of roughly 150 per pixel.," A total of 38 spectra were taken covering the interval, with estimated signal to noise ratios of roughly 150 per pixel."1144 One spectrum was taken coveringAA.. aimed at the line.," One spectrum was taken covering, aimed at the $\beta$ line."1145 The reduction of all spectra was undertaken using optimal extraction (Horne1986) as implemented in the software tools and Marsh¢1989)., The reduction of all spectra was undertaken using optimal extraction \citep{Horne86pasp} as implemented in the software tools and \citet{Marsh89pasp}.1146", Our main photometric dataset comes from the WASP survey for transiting extrasolar planets (Pollaccoetal.2006) which operates two instruments. one in La Palma and one in South Africa."," Our main photometric dataset comes from the WASP survey for transiting extrasolar planets \citep{Pollacco+06pasp} which operates two instruments, one in La Palma and one in South Africa."1147 Each instrument consists of eight small telescopes on a common mount. and each telescope comprises a Canon mmm telephoto lens. a custom resembling g|r (Fukugitaetal.1996).. and an E2V 2k CCD.," Each instrument consists of eight small telescopes on a common mount, and each telescope comprises a Canon mm telephoto lens, a custom resembling $g$$+$$r$ \citep{Fukugita+96aj}, and an E2V $\times$ 2k CCD."1148 Both instruments studied the area surrounding CCet extensively in the 2008 and 2009 observing seasons., Both instruments studied the area surrounding Cet extensively in the 2008 and 2009 observing seasons.1149 The data were reduced by a dedicated pipeline (Pollaccoetal.2006) which performs aperture photometry in three software apertures and then applies the detrending algorithm (Tamuz to the photometry from the second of these apertures., The data were reduced by a dedicated pipeline \citep{Pollacco+06pasp} which performs aperture photometry in three software apertures and then applies the detrending algorithm \citep{Tamuz++05mn} to the photometry from the second of these apertures.1150Usiug an initial seed based on carly Ajσ relations produces broad scatter in the final relations.,Using an initial seed based on early $\msig$ relations produces broad scatter in the final relations.1151" This is especially evident in the ""ni-s.halo.log"" combination. m which there appear to be two distinct. populations of SMDIIs: one population aloug the observed relation aud another at lower mass."," This is especially evident in the “m-s,halo,log” combination, in which there appear to be two distinct populations of SMBHs: one population along the observed relation and another at lower mass."1152 This behavior is due to the fact hat halos are seeded in two scenarios: when the halo first becomes resolvable iu the simulation. and if the ialo merges with another halo. loses its black hole. aud ater separates.," This behavior is due to the fact that halos are seeded in two scenarios: when the halo first becomes resolvable in the simulation, and if the halo merges with another halo, loses its black hole, and later separates."1153 Figure 10. illustrates this bv showiug he initial seed mass as a function of redshift., Figure \ref{fig:bhinit} illustrates this by showing the initial seed mass as a function of redshift.1154 Note that seed black holes are never larger than ~10.7? of the host ialo mass., Note that seed black holes are never larger than $\sim 10^{-5}$ of the host halo mass.1155 Small halos coutiuue to appear throughout he evolution of the simulation and are seeded with ~10+AL. black holes., Small halos continue to appear throughout the evolution of the simulation and are seeded with $\sim 10^4 \msol$ black holes.1156 However. occasionally a lavecr ialo loses its black hole and nmst be re-seeded with a correspondingly larger SMDIT.," However, occasionally a larger halo loses its black hole and must be re-seeded with a correspondingly larger SMBH."1157" Thus we may be left with wo popluations: halos with their original SMDIIS. that evolve to relations simular to halos with coustaut iitial MASS, al Lhalos that are re-seeded at lower redshifts with lack. ho Classes closer to the observed relation."," Thus we may be left with two populations: halos with their original SMBHs that evolve to relations similar to halos with constant initial mass, and halos that are re-seeded at lower redshifts with black hole masses closer to the observed relation."1158 This (listiuctikn is largely eliminated by applviug the merger ests at every time step aud mereie black holes iustautly ou halo mergers., This distinction is largely eliminated by applying the merger tests at every time step and merging black holes instantly on halo mergers.1159 This results both in higher mass due Oo increased ποσο and in increased low-redshift το-seeding of high-mass halos., This results both in higher mass due to increased mergers and in increased low-redshift re-seeding of high-mass halos.1160 This implies that re-seeding selects a few high-mass halos aud places them on the observed relation at late times., This implies that re-seeding selects a few high-mass halos and places them on the observed relation at late times.1161 Indeed. we find that all," Indeed, we find that all"1162where the factor (dI/dw) describes the intensity of the radiation emitted by the electron in nultiple collisions.,where the factor $\left\langle dI/d\omega \right\rangle$ describes the intensity of the radiation emitted by the electron in multiple collisions.1163 We assume that the electrons form a degenerate Fermi gas. wilh (he particle number density 7 given by (Chin1963) The emissivity per unit volume ol the electrosphere can be calculated by using the delinilion ↥≀↧∶∖↽≼↲↕⋅∪↓≯⊔∐≺∢↳↽∐≼↲⊳∖⋱∖⊽↙∣⊳∶⋅↥⋟∖⇁⊼↼⋛↙⊔∶∶↸−↴⊔ ⋣↙∣⊳∶∕∕∕⊼⋅↴⊺≀↧↴↳↽↥∐≸≟↕∐↥∪≀↧↴≺∢≺∢∪∏∐↥⊔∐↲≺∢∪∐∏⋅↕∣↽≻∏∐∪∐∪↓⋟≀↧↴∐↥≀↧∶∖↽≼↲↕⋅⊳∖⇁ ve find 1]," We assume that the electrons form a degenerate Fermi gas, with the particle number density $n$ given by \citep{Ch68}1164 The emissivity per unit volume of the electrosphere can be calculated by using the definition The energy flux from the electrosphere of a strange star, coming out from a thin surface layer of thickness $dz$, is $F_{\gamma1165}^{(ee)}=\varepsilon _{\gamma }^{(ee)}dz/\pi $."1166 As a first step in obtaining the photon emissivity of the electrosphere of the bare strange stars we have to obtain the classical radiation intensity distribution <df/duo>. emitted by electrons moving in a dense medium in which many inter-parlicle collisions occur.," Taking into account the contribution of all layers we find As a first step in obtaining the photon emissivity of the electrosphere of the bare strange stars we have to obtain the classical radiation intensity distribution $<dI/d\omega >$, emitted by electrons moving in a dense medium in which many inter-particle collisions occur."1167 Since in the electrosphere a very strong electric field is also present. the influence of this field on the raciation process must also be taken into account.," Since in the electrosphere a very strong electric field is also present, the influence of this field on the radiation process must also be taken into account."1168 In (he early 1950s Landau&Pomeranchuk(1953) and Mieclal (1956)ppredicted that the racialion from highlv relativistic particles in dense meclia is suppressed due to interference between amplitudes of nearby interactions (for detailed presentations of the LPAI effect see ihe monograph by AkhiezerandShul'gea(1996) and the review article by Ixlein(1999)))., In the early 1950s \citet{LaPo53} and \citet{Mi56}p predicted that the radiation from highly relativistic particles in dense media is suppressed due to interference between amplitudes of nearby interactions (for detailed presentations of the LPM effect see the monograph by \citet{AkSh96} and the review article by \citet{Kl99}) ).1169 The suppression effect is of quantum mechanic nature and has its roots in the uncertainty principle., The suppression effect is of quantum mechanic nature and has its roots in the uncertainty principle.1170 The kinematics of the bremsstrahlune requires that the longitudinal momentum transfer between (he interacting particles must be small: (he uncertainty principle then requires (hat (he interaction must occur over a large longitudinal distance scale (formation zone)., The kinematics of the bremsstrahlung requires that the longitudinal momentum transfer between the interacting particles must be small; the uncertainty principle then requires that the interaction must occur over a large longitudinal distance scale (formation zone).1171 If the electron Coulomb scatters while traversing (is zone. (he bremsstralilung aanplitude from before aud after scattering can interfere. (hus reducing (he amplitude for bremsstralilung photon emission (AkhiezerandShul'ga(1996) and references therein).," If the electron Coulomb scatters while traversing this zone, the bremsstrahlung amplitude from before and after scattering can interfere, thus reducing the amplitude for bremsstrahlung photon emission \citet{AkSh96} and references therein)."1172 The results of the quantitative measurements of the bremsstrahlung suppression due to the LPM effect. [or, The results of the quantitative measurements of the bremsstrahlung suppression due to the LPM effect for1173"the ""hot spot” owing to. e.g.. rapid rotation (x long lasting acerction flow.","the “hot spot” owing to, e.g., rapid rotation or long lasting accretion flow."1174" It is conceivable tjw onlv a limited fraction of superuovae create the ""host spot relevant for the current scenario (0.8... rapid rotators or less-energetie supernovae tjit. form(nmm the5 longo lastinglastine accretiaccretion flow)."," It is conceivable that only a limited fraction of supernovae create the “host spot” relevant for the current scenario (e.g., rapid rotators or less-energetic supernovae that form the long lasting accretion flow)."1175M Thislis can be a reasonable explanation for that the spectroscopic analysis of extremely imetal-poor stars and Galactic chemical evolution study imply ouly a Iunited fraction of core-collapse superunovae undergo the r-process nucleosyuthesis 1)., This can be a reasonable explanation for that the spectroscopic analysis of extremely metal-poor stars and Galactic chemical evolution study imply only a limited fraction of core-collapse supernovae undergo the -process nucleosynthesis .1176. The implications in this study iust be tested by future multi-dinmeusional simulations of core-collapse superuovae for long duration (~ 105) with accurate neutrino transport.," The implications in this study must be tested by future multi-dimensional simulations of core-collapse supernovae for long duration $\sim 10\, \mathrm{s}$ ) with accurate neutrino transport."1177 Systematic calculationsof nucleosvuthesis with such hivdrodyviiuauic trajectoriesJ will be also needed to investigateD the contribution to the GalacticUt chemical| evolution of r--process nuclei., Systematic calculations of nucleosynthesis with such hydrodynamic trajectories will be also needed to investigate the contribution to the Galactic chemical evolution of -process nuclei.1178 I would like to acknowledge IL. -Th., I would like to acknowledge H. -Th.1179 Janka for helpful discussions., Janka for helpful discussions.1180 I also acknowledge the contributions of ai anouyviuous referee. which led to claxification ofa uunuber of poiuts in the original manuscript.," I also acknowledge the contributions of an anonymous referee, which led to clarification of a number of points in the original manuscript."1181 This work was supported iu part bv a Carant-in-Aid for Scicutifie Research (17710108) frou. the Ministry of Education. Culture. Sports. Science. and Technology of Japan.," This work was supported in part by a Grant-in-Aid for Scientific Research (17740108) from the Ministry of Education, Culture, Sports, Science, and Technology of Japan."11822010).,.1183.. While including a more complicated biasing model may improve agreement with the pairwise velocity statistics of halos. our results suggest that these corrections are small on the scales of interest here.," While including a more complicated biasing model may improve agreement with the pairwise velocity statistics of halos, our results suggest that these corrections are small on the scales of interest here."1184 Analogous to Eq. 27..," Analogous to Eq. \ref{vinfallPT},"1185 we compute the pair-weighted velocity dispersion., we compute the pair-weighted velocity dispersion.1186" The result depends only on ge= cos*(dy,.). where ó;, is the angle between the LOS and the pair separation vector."," The result depends only on $\mu^2 = \cos^2(\phi_{\ell r})$ , where $\phi_{\ell r}$ is the angle between the LOS and the pair separation vector."1187 As in Eq. |l..," As in Eq. \ref{vdisptot},"1188 we present results for the velocity dispersion perpendicular and parallel to the LOS. which can be combined to give the dispersion as a function of Gr.£CY.," we present results for the velocity dispersion perpendicular and parallel to the LOS, which can be combined to give the dispersion as a function of $(r, \mu^2)$."1189 We again separate the terms by their bias and. scale dependence. and provide explicit expressions in Appendix Az: The higher order terms in line 29. ean be accounted for in he form of Equations 9. and. 10... but replacing ΟΕΚ} with the xerturbation theory result for P..(A).," We again separate the terms by their bias and scale dependence, and provide explicit expressions in Appendix \ref{ptcalcs}: The higher order terms in line \ref{v2r1} can be accounted for in the form of Equations \ref{psiperp} and \ref{psiparallel}, , but replacing $P_m^r(k)$ with the perturbation theory result for $P_{\theta \theta}(k)$."1190 The term in line 30. evaluates ο a constant. which for our fiducial cosmological parameters is 13.245/7. Uh'Mper.," The term in line \ref{v2r2} evaluates to a constant, which for our fiducial cosmological parameters is $13.24b f^2$ $h^{-1}$ ${}^2$."1191 While the perturbation theory calculation overestimates the amplitude of the effect. Fig.," While the perturbation theory calculation overestimates the amplitude of the effect, Fig."1192 9 does show a arge scale offset between the linear theory velocity dispersions and those measured for halos in our simulations: the offset for our HOD sample is ~1.501'Mpey.," \ref{fig:vstats} does show a large scale offset between the linear theory velocity dispersions and those measured for halos in our simulations; the offset for our HOD sample is $\sim 1.5\,(h^{-1}{\rm Mpc})^2$."1193 There is a slight dependence on ilo mass. with the high mass sample offset ~1.007'Mper and he low mass sample offset ~2.2Ur'Mpey.," There is a slight dependence on halo mass, with the high mass sample offset $\sim 1.0\,(h^{-1}{\rm Mpc})^2$ and the low mass sample offset $\sim 2.2\,(h^{-1}{\rm Mpc})^2$ ."1194 The offsets for the xrallel and perpendicular components are in agreement with each other at the level of 0.1(7'Mpey.," The offsets for the parallel and perpendicular components are in agreement with each other at the level of $0.1\,(h^{-1}{\rm Mpc})^2$."1195 In a future paper we expect to accommodate this offset into our theory along with other small-scale isotropic dispersions due to redshift errors and fingers-of-god from satellite galaxies., In a future paper we expect to accommodate this offset into our theory along with other small-scale isotropic dispersions due to redshift errors and fingers-of-god from satellite galaxies.1196 However. wepoint out that for halos. the small-scale. incoherent velocity dispersion is much smaller than the linear theory dispersion o;=210:'Mper that appears in Eq. 24..," However, wepoint out that for halos, the small-scale, incoherent velocity dispersion is much smaller than the linear theory dispersion $\sigma_v^2 = 21\,(h^{-1}{\rm Mpc})^2$ that appears in Eq. \ref{scoccPT}."1197 Line 31 has many terms that depend on By. we evaluate them explicitly in the Appendix A.., Line \ref{v2r4} has many terms that depend on $B_{\delta \theta \theta}$; we evaluate them explicitly in the Appendix \ref{ptcalcs}.1198 Finally. line 32. reduces to This shows that the pair-weighting factor cancels out to leading order. and the dispersion should be increased along the separation vector due to the linear infall.," Finally, line \ref{v2r6} reduces to This shows that the pair-weighting factor cancels out to leading order, and the dispersion should be increased along the separation vector due to the linear infall."1199 The middle and right panels of Fig., The middle and right panels of Fig.1200 9. compare the scale dependence of the HOD halo subsample velocity dispersions parallel and perpendicular to the LOS. respectively.," \ref{fig:vstats} compare the scale dependence of the HOD halo subsample velocity dispersions parallel and perpendicular to the LOS, respectively."

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