ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 Of all the single-satellite data sets that we consider. the IIETE data is the only one that apparently obevs the Amati relation. although its large scatter limits ils usefulness.," Of all the single-satellite data sets that we consider, the HETE data is the only one that apparently obeys the Amati relation, although its large scatter limits its usefulness."3 For the data. we use the catalog in Butler οἱ al. (," For the data, we use the catalog in Butler et al. ("42007).,2007).5" We use their Z4, as derived from frequentist statistics as it is (he most common approach to finding Ej. ("," We use their $E_{peak,obs}$ as derived from frequentist statistics as it is the most common approach to finding $E_{peak}$. ("6Their Bavesian values were mace with unreasonable priors that signilicantlv skew (he results.),Their Bayesian values were made with unreasonable priors that significantly skew the results.)7 The burst detector only goes up to. 150 keV. so the reported values of E;4. are almost all lower Chan 200 keV. We have adopted their bolometric fluences aud we have converted (heir non-standard error bars into standard one-sigma error bars.," The burst detector only goes up to 150 keV, so the reported values of $E_{peak,obs}$ are almost all lower than 200 keV. We have adopted their bolometric fluences and we have converted their non-standard error bars into standard one-sigma error bars."8 Figure 9. plots the results: bursts without. known redshilts are represented as empty circles and bursts with known redshift are represented by a filled diamond., Figure \ref{fig:NaPSwift} plots the results; bursts without known redshifts are represented as empty circles and bursts with known redshift are represented by a filled diamond.9 This will be true for all future plots., This will be true for all future plots.10" This is the last of the three plots in which we have plotted lines representing (rigeer (dashed line) and E,,,; detection thresholds (dot-dashed line). which again. are detailed in section 4."," This is the last of the three plots in which we have plotted lines representing trigger (dashed line) and $E_{peak}$ detection thresholds (dot-dashed line), which again, are detailed in section 4."11 The bursts violate the Amati limit at a rate of to824., The bursts violate the Amati limit at a rate of to.12 That is. the Amati relation does not work [orSiwift.," That is, the Amati relation does not work for."13. This result is not the result of small number statistics. and we ean see from the distribution that the disagreement is highly significant.," This result is not the result of small number statistics, and we can see from the distribution that the disagreement is highly significant."14 This is another version of the same conclusion first reported by Butler et al. (, This is another version of the same conclusion first reported by Butler et al. (152007).,2007).16wavelengths using single dish telescopes (Ixastner οἱ al.,wavelengths using single dish telescopes (Kastner et al.17 1997. van Zadelhoff et al.," 1997, van Zadelhoff et al."18 2001)., 2001).19 More recent. deeper. observations have also detected emission from LIMCO and . which indicate deuterium fractionation in the disk similar to cold cores and pristine cometary malerial (van Dishoeck οἱ al.," More recent, deeper, observations have also detected emission from $^{13}$ $^{+}$ and $^{+}$, which indicate deuterium fractionation in the disk similar to cold cores and pristine cometary material (van Dishoeck et al."20 2003)., 2003).21 So lar. no observations have been made that spatially resolve the molecular gas surrounding TW Ilva.," So far, no observations have been made that spatially resolve the molecular gas surrounding TW Hya."22 Such spatially resolved. observations are important to verily underlving constructs of the models used to interpret the single dish data. since the physical conditions are subject to strong radial and vertical gradients. which result in considerable chemical complexity.," Such spatially resolved observations are important to verify underlying constructs of the models used to interpret the single dish data, since the physical conditions are subject to strong radial and vertical gradients, which result in considerable chemical complexity."23 IID 100546. among the nearest IIerbig Ae/Be stars at 10346 pe (Lipparcos). appears nearby the dark cloud. DC296.2-7.9 αμα. The de Winter 1989) and shows a disk-like scaltered light distribution (Pantin. Waelkens Lagage 2000. Augereau οἱ al.," HD 100546, among the nearest Herbig Ae/Be stars at $103\pm6$ pc (Hipparcos), appears nearby the dark cloud DC296.2-7.9 (Hu, The de Winter 1989) and shows a disk-like scattered light distribution (Pantin, Waelkens Lagage 2000, Augereau et al."24 2001. Grady el al.," 2001, Grady et al."25 2001) of substantial size. about 8” (8000 AU). as well as strong millimeter emission from dust (Ilenning et al.," 2001) of substantial size, about $8''$ (8000 AU), as well as strong millimeter emission from dust (Henning et al."26 1994. 1993).," 1994, 1998)."27 This relatively isolated svstem is thought to have an age of ~10 Myr (van den Ancker et al., This relatively isolated system is thought to have an age of $\sim10$ Myr (van den Ancker et al.28 1997)., 1997).29 However. analvsis of the full spectral enerev distribution suggests the presence of an extended. envelope (Ilenning et al.," However, analysis of the full spectral energy distribution suggests the presence of an extended envelope (Henning et al."30 1994)., 1994).31 Alid-infrared. spectroscopy from ISO shows remarkably strong crvstalline silicate. bands. similar to those observed in comet Hale-Dopp (Malfait et al.," Mid-infrared spectroscopy from ISO shows remarkably strong crystalline silicate bands, similar to those observed in comet Hale-Bopp (Malfait et al."32 1998) and indicative of unusual and substantial processing of the dust. within the disk (Bouwman et al., 1998) and indicative of unusual and substantial processing of the dust within the disk (Bouwman et al.33 2003)., 2003).34 Millimeter interferometry has (he potential to show directly the presence of a disk component that contains the bulk of the svstem mass., Millimeter interferometry has the potential to show directly the presence of a disk component that contains the bulk of the system mass.35 We observed TW Ilva and WD 100456 at 89 GIIz (3.4 mm) with the ATCA during the 2002 austral winter using (wo compact configurations of three 22 meter diameter antennas equipped with millimeter receivers., We observed TW Hya and HD 100456 at 89 GHz (3.4 mm) with the ATCA during the 2002 austral winter using two compact configurations of three 22 meter diameter antennas equipped with millimeter receivers.36 Table 1 summarizes (he observational parameters., Table 1 summarizes the observational parameters.37" The ATCA observations provided 6 independent east-west baselines with lengths ranging [rom 22 to 230 meters. resulting in ~2"" resolution."," The ATCA observations provided 6 independent east-west baselines with lengths ranging from 22 to 230 meters, resulting in $\sim2''$ resolution."38" For each observation. the target pointing center was set to be 5"" west of the star position."," For each observation, the target pointing center was set to be $5''$ west of the star position."39 The digital correlator was configured for two dual polarization bands: (1) a wide band with 33 channels spanning the maximum bandwidth of 128 MIIz for continuum sensitivity. and (2) a narrow band with 256 channels over 16 MlIz to provide high frequency resolution on the J=10 line. which is one of the few potentially detectable spectral lines accessible to the interim svstem.," The digital correlator was configured for two dual polarization bands: (1) a wide band with 33 channels spanning the maximum bandwidth of 128 MHz for continuum sensitivity, and (2) a narrow band with 256 channels over 16 MHz to provide high frequency resolution on the $^+$ J=1–0 line, which is one of the few potentially detectable spectral lines accessible to the interim system."40 The sky signal is mixed with a Gunn oscillator locked at 80505.5 MIIz ancl passed through an (58000-10800. MIIZ) filter aud splitter module. which results in an accessible frequency range of 88.506 to 91.305 GlIz.," The sky signal is mixed with a Gunn oscillator locked at 80505.5 MHz and passed through an X-band (8000-10800 MHz) filter and splitter module, which results in an accessible frequency range of 88.506 to 91.305 GHz."41 A new local oscillator system under construction will allow, A new local oscillator system under construction will allow42We used archive imaging data of 3346 110248. ΡΙ. A. Nota).,"We used archive imaging data of 346 10248, P.I. A. Nota)."43 These observations were obtained with ACS (Advanced Camera for Surveys) onboardHST., These observations were obtained with ACS (Advanced Camera for Surveys) onboard.44 The images were taken with the Wide Field Camera (WFC) using broad- and narrow-band filters (FS55W. F814W. F656N) in 2004 July.," The images were taken with the Wide Field Camera (WFC) using broad- and narrow-band filters (F555W, F814W, F656N) in 2004 July."45 They were used to examine the morphology of the compact rregion N66À and in. particular resolve its exciting stars., They were used to examine the morphology of the compact region N66A and in particular resolve its exciting stars.46 In addition. we produced a composite image of N66A . 3)undalsousedthe photometryderived fromthe seboMdrvations(Goulierntis2006).," In addition, we produced a composite image of N66A \\ref{fig:HST_neb}, \ref{fig:HST_stars}) ) and also used the photometry derived from these observations (Gouliermis 2006)."47 Weal sousedthearchiveob servationsobtai ," We also used the archive observations obtained in July 2006 using ACS with the High Resolution Channel (HRC) and the ultraviolet filters F220W and F330W 10542, P.I. A. Nota)."48The Spitzer archive data used in this paper come from the S*MC project., The Spitzer archive data used in this paper come from the $^{3}$ MC project.49 This is a project to map the star-forming body of the SMC with Spitzer in all seven Infrared Array Camera (IRAC) and Multiband Imaging Photometer for Spitzer (MIPS) bands., This is a project to map the star-forming body of the SMC with Spitzer in all seven Infrared Array Camera (IRAC) and Multiband Imaging Photometer for Spitzer (MIPS) bands.50 The MIPS data were obtained in 2004 November and the IRAC data in 2005 May., The MIPS data were obtained in 2004 November and the IRAC data in 2005 May.51 We used the IRAC data to build a composite image of N66 and also obtain the photometry of NOOA. The typical PSF of the IRAC images in the 3.6. 4.5. 5.8. and 8.0 jm bands ts [7.06 to 1.98 and that of MIPS at 24 pm is (Bolattoetal..2007).," We used the IRAC data to build a composite image of N66 and also obtain the photometry of N66A. The typical PSF of the IRAC images in the 3.6, 4.5, 5.8, and 8.0 $\mu$ m bands is .66 to .98 and that of MIPS at 24 $\mu$ m is \citep[][]{Bolatto07}."52. The derived photometry for N66A in the 3.6. 4.5. 5.8. and 8.0 uim bands are 12.22. 11.82. 10.25. and 8.66 mag. respectively. using an integration aperture of 5 pixels. or in radius.," The derived photometry for N66A in the 3.6, 4.5, 5.8, and 8.0 $\mu$ m bands are 12.22, 11.82, 10.25, and 8.66 mag, respectively, using an integration aperture of 5 pixels, or in radius."53 Measurements with either slightly larger or smaller apertures do not affect the color results., Measurements with either slightly larger or smaller apertures do not affect the color results.54 Moreover. we tried to detect substructures in the Spitzer images of N66À (see Sect.," Moreover, we tried to detect substructures in the Spitzer images of N66A (see Sect."55 4)., 4).56 Figure | displays a composite image of the N66 region taken with the NTT telescope (see Sect., Figure 1 displays a composite image of the N66 region taken with the NTT telescope (see Sect.57 2.1)., 2.1).58 The field is corresponding to «000 pe? for a distance of about 60 kpe. orM = 18.94 mag (Laney&Stobie.1994).," The field is corresponding to 90 $^{2}$ for a distance of about 60 kpc, or = 18.94 mag \citep[][]{Laney94}."59. Although the resolution ts less than that of the ACS image (Sect., Although the resolution is less than that of the ACS image (Sect.60 2.3.1). almost all the features of the ccomplex are visible.," 2.3.1), almost all the features of the complex are visible."61 The 3346 cluster appears to be at the center of an bbowl. the southern border of which is delineated by a compressed tonized gas front and an absorption lane running over some 60 pc.," The 346 cluster appears to be at the center of an bowl, the southern border of which is delineated by a compressed ionized gas front and an absorption lane running over some 60 pc."62 In particular. N66A stands out as the most compact nnebula of the whole region. with coordinates (J2000.0) a = 00:59:14.8. 0 = -72:11:01.," In particular, N66A stands out as the most compact nebula of the whole region, with coordinates (J2000.0) $\alpha$ = 00:59:14.8, $\delta$ = -72:11:01."63 The compact rregion is apparently associated with the compressed gas front and the absorption lane., The compact region is apparently associated with the compressed gas front and the absorption lane.64 The field of view of the NTT image is larger than that ofHST ACS., The field of view of the NTT image is larger than that of ACS.65 It also displays a turbulent environment in the eastern side of N66 with many indications of shocked gas., It also displays a turbulent environment in the eastern side of N66 with many indications of shocked gas.66 In particular. the wind-driven bubble centered on 55980 is quite impressive.," In particular, the wind-driven bubble centered on 5980 is quite impressive."67 A narrow ridge can also be discerned towards the southern outer boundary of the complex., A narrow ridge can also be discerned towards the southern outer boundary of the complex.68 This feature isalso affected by stellar shock winds. as indicated by its remarkable eemission (Reidetal..2006).," This feature isalso affected by stellar shock winds, as indicated by its remarkable emission \citep[][]{Reid06}."69. Observations with the Australia Telescope Compact Array (ATCA) and Parkes Observatory at the Australia Telescope National Facility (ATNF) also detect an ccloud concentrated towards this part of N66 (Staveley-Smithetal..1997;Stanimirovié1999:Gouliermis2008 ).," Observations with the Australia Telescope Compact Array (ATCA) and Parkes Observatory at the Australia Telescope National Facility (ATNF) also detect an cloud concentrated towards this part of N66 \citep[][]{Staveley-Smith97,Stanimirovic99,Gouliermis08}."70. The ccloud is probably 1n contact with the ridge., The cloud is probably in contact with the ridge.71 Figure 2 presents a high-resolution composite image of the N66A rregion extracted from the above-mentionedHST ACS archive data., Figure \ref{fig:HST_neb} presents a high-resolution composite image of the N66A region extracted from the above-mentioned ACS archive data.72 The compact rregion is in diameter. corresponding to .33 pe.," The compact region is in diameter, corresponding to 3 pc."73 It contains a strong absorption lane., It contains a strong absorption lane.74 Interestingly. two bright stars. labelled #11 and #22. are located towards the central part of the region. above the dust lane (see also reftig:HST .tars).," Interestingly, two bright stars, labelled 1 and 2, are located towards the central part of the region, above the dust lane (see also \\ref{fig:HST_stars}) )."75SeparatedbyQ 7 .00.2 pe). they are the main exciting stars of the rregion. as shown in Sect.," Separated by .7 0.2 pc), they are the main exciting stars of the region, as shown in Sect."76 3.3., 3.3.77 A number of fainter stars are seen across the face of the region., A number of fainter stars are seen across the face of the region.78 These stars are quite bright on theHST ACS/HRC image in the ultraviolet obtained with F220W and F330W filters., These stars are quite bright on the ACS/HRC image in the ultraviolet obtained with F220W and F330W filters.79 Five other stars in the UV images (#33 to #77) should also be OB stars associated with N66AÀ. as suggested by their photometry. presented in Table 2.. which also displays the positions.," Five other stars in the UV images 3 to 7) should also be OB stars associated with N66A, as suggested by their photometry, presented in Table \ref{table:A}, which also displays the positions."80 The total H8 flux of N66A was derived using the following procedure., The total $\beta$ flux of N66A was derived using the following procedure.81 First we calculated the relative HB flux in ar imaginary slit passing through the Hg image with respect to the total flux emitted by the whole NO6A region., First we calculated the relative $\beta$ flux in an imaginary slit passing through the $\beta$ image with respect to the total flux emitted by the whole N66A region.82 This value was ther compared with the absolute flux obtained from the spectra., This value was then compared with the absolute flux obtained from the spectra.83 In both cases. a mean flux measured for the NS and EW orientations of the slit was used.," In both cases, a mean flux measured for the NS and EW orientations of the slit was used."84 The total Hf flux thus obtainec was F(Hp) erg em s7!., The total $\beta$ flux thus obtained was $F$ $^{-12}$ erg $^{-2}$ $^{-1}$.85 Considering the extinction law for the LMC (Howarth. 1983).. we computed the reddening corrected intensity /(H8) ," Considering the extinction law for the LMC \citep{Howarth83}, , we computed the reddening corrected intensity $I$ "86teiiperature code using fus-linited radiative diffusion.,temperature code using flux-limited radiative diffusion.87 It functions very well. however. in situations where the light curve is recombination donimnated (as on the plateau of Type ILP superuovae). in the radioactive peaks of Type I supernova. and in the tails of Type IEP supernovae.," It functions very well, however, in situations where the light curve is recombination dominated (as on the plateau of Type II-P supernovae), in the radioactive peaks of Type I supernova, and in the tails of Type II-P supernovae."88" CGanmuna-deposition has been calibrated against a Monte Carlo code (Pinto and Woosley 1988) aud the code was successfully used to predict the behavior of SN. 1987À (Woosley, Pinto. Eusuun 1988). though muxing was later added as an esseutial ineredicut."," Gamma-deposition has been calibrated against a Monte Carlo code (Pinto and Woosley 1988) and the code was successfully used to predict the behavior of SN 1987A (Woosley, Pinto, Ensman 1988), though mixing was later added as an essential ingredient."89 Because thermal equilibrium and black body spectra are assumed. the bolometric luminosity is known much better than the color photometry.," Because thermal equilibrium and black body spectra are assumed, the bolometric luminosity is known much better than the color photometry."90 Thus D and V imaguitudes should be approximately correct. but U magnitudes are much more Uncertain.," Thus B and V magnitudes should be approximately correct, but U magnitudes are much more uncertain."91 The resulting luminosities. cffective teniperatures. shotospheric radii are shown in Figure d1..," The resulting luminosities, effective temperatures, photospheric radii are shown in Figure \ref{fig:sne}."92— As the shock uoves toward the low-densityv stellar surface. its energv is deposited iuto progressively smaller amounts of matter.," As the shock moves toward the low-density stellar surface, its energy is deposited into progressively smaller amounts of matter."93 This results in high velocities aud temperatures when he shock reaches the stellar surface. causing a pulse of ultraviolet radiation with a characteristic timescale of a ew nutes.," This results in high velocities and temperatures when the shock reaches the stellar surface, causing a pulse of ultraviolet radiation with a characteristic timescale of a few minutes."94 This “breakout” phase is by far tle most 1unuinous and bluest phase of the PISN burst. but its very short duration makes it dificult to use in observational searches.," This “breakout” phase is by far the most luminous and bluest phase of the PISN burst, but its very short duration makes it difficult to use in observational searches."95 Iu fact. the analog of this pliase in couvetional SN has so far only beeu iuivectly detected iu SN 1987AÀ (Navozliu 199I: Ibuuuy 1988: Catchpole 1055).," In fact, the analog of this phase in conventional SN has so far only been indirectly detected in SN 1987A (Nayozhin 1994; Hamuy 1988; Catchpole 1988)."96" Followiug breakout. the st: woexpanuds with f, 1utially proportional to tine."," Following breakout, the star expands with $R_{\rm phot}$ initially proportional to time."97 Though a simall fraction of the outer dnass may move much faster. the characteristic velocity of the photosphere during this pliase is a modest eo=(Q2NE/M)/?~(OMeres/200AE.82—5000 lin/s. because of the very laree mass participating in the explosion.," Though a small fraction of the outer mass may move much faster, the characteristic velocity of the photosphere during this phase is a modest $v = (2 KE/M)^{1/2} \sim (10^{53} {\rm98ergs}/200 \msun)^{1/2} \sim 5000$ km/s, because of the very large mass participating in the explosion."99" During the expansion. the radiatiou- ejecta cool adiabatically. with Z approximately proportional to R1, with au additional cucrey input from the decay of °ONi (f a significant mass was svuthesized during the explosion) aud hydrogen recombinations Gvleu T~ 104)."," During the expansion, the radiation-dominated ejecta cool adiabatically, with $T$ approximately proportional to $R^{-1},$ with an additional energy input from the decay of $^{56}$ Ni (if a significant mass was synthesized during the explosion) and hydrogen recombinations (when $T \sim 10^4$ K)."100 As the scale radius for this cooling is the radius of the progenitor. the temperatures and bpuninosities are substantially larger throughout this phase in the cases witli the strongest mixiug.," As the scale radius for this cooling is the radius of the progenitor, the temperatures and luminosities are substantially larger throughout this phase in the cases with the strongest mixing."101 After ~50 davs. the οπσον input from Co decay becomes larger than the remaining thermal cnerev (the initial thermal enerev deposited bv the shock is mostly eaten away by the adiabatic expausion).," After $\sim 50$ days, the energy input from $^{56}$ Co decay becomes larger than the remaining thermal energy (the initial thermal energy deposited by the shock is mostly eaten away by the adiabatic expansion)."102" At this time the cnerey deposited by ""Co in deeper lavers that were", At this time the energy deposited by $^{56}$ Co in deeper layers that were10310002000kins+). M~10.1AL. yr1) has been achieved. mostly due to low statistics and to the relatively poor spectral resolution of the detectors.,"$v_w\sim 1000-2000\kms$ $\mdot\sim 10^{-5}-10^{-4}\msole$ $^{-1}$ has been achieved, mostly due to low statistics and to the relatively poor spectral resolution of the detectors."104 Here. we consider the Swift (Gehrels et al.," Here, we consider the Swift (Gehrels et al."105 2004) observations of GRBO060218., 2004) observations of GRB060218.106 This ts the second closest GRB (redshift 0.033). and the first one showing the shock break out of the SN (Campana et al.," This is the second closest GRB (redshift $z=0.033$ ), and the first one showing the shock break out of the SN (Campana et al."107 2006)., 2006).108 Modelling of the spectra and light curve of the associated SN 2006aj (Pian et al., Modelling of the spectra and light curve of the associated SN 2006aj (Pian et al.109 2006: Mirabal et al., 2006; Mirabal et al.110 2006: Sollerman et al., 2006; Sollerman et al.111 2006: Cobb et al., 2006; Cobb et al.112 2006) suggested a progenitor star whose initial mass was 20+1AL. (Mazzali et al., 2006) suggested a progenitor star whose initial mass was $20\pm1 \msole$ (Mazzali et al.113 2000)., 2006).114 This GRB was of very long duration. which allowed Swift to observe it with its narrow field instruments (the X-ray Telescope. XRT. Burrows et al.," This GRB was of very long duration, which allowed Swift to observe it with its narrow field instruments (the X–ray Telescope, XRT, Burrows et al."115 2005. and the UV/Optical Telescope. UVOT) during a considerable part of its prompt phase. collecting the largest number of X-ray photons ever.," 2005, and the UV/Optical Telescope, UVOT) during a considerable part of its prompt phase, collecting the largest number of X–ray photons ever."116 In this letter we exploit the huge number of X-ray photons by analysing the absorption pattern burnt into it by cireumburst material., In this letter we exploit the huge number of X–ray photons by analysing the absorption pattern burnt into it by circumburst material.117 The XRT spectra have been obtained in the Windowed Timing (WT) mode in which a ID image is obtained by adding the data along the central 200 pixels in a single row (see Hill et al., The XRT spectra have been obtained in the Windowed Timing (WT) mode in which a 1D image is obtained by adding the data along the central 200 pixels in a single row (see Hill et al.118 2004)., 2004).119 The XRT data have been processed using the FTOOLS software package (v. 6.3.1) distributed within HEASOFT., The XRT data have been processed using the FTOOLS software package (v. 6.3.1) distributed within HEASOFT.120 We run the taskartpipeline (v.0.11.4) applying calibration and standard filtering and screening criteria., We run the task (v.0.11.4) applying calibration and standard filtering and screening criteria.121 In particular. we dynamically correct for possible bias offsets computing the bias difference between the on-ground estimated bias median from the last 20 pixels data telemetered with every frame. and the median of the last 20 pixels in the related bias row.," In particular, we dynamically correct for possible bias offsets computing the bias difference between the on-ground estimated bias median from the last 20 pixels data telemetered with every frame, and the median of the last 20 pixels in the related bias row."122 Events with grade 0 have been selected in order to attain the best spectral resolution., Events with grade 0 have been selected in order to attain the best spectral resolution.123 The XRT analysis has been performed in the 0.3-10 keV energy band (and also in the 0.35-10 keV and 0.25-10 keV energy bands as a consistency check)., The XRT analysis has been performed in the 0.3–10 keV energy band (and also in the 0.35–10 keV and 0.25–10 keV energy bands as a consistency check).124 Given the XRT CCD resolution at low energies (AL/E~ 15%) we cannot directly see the edges imprint in the X-ray spectrum., Given the XRT CCD resolution at low energies $\Delta E/E\sim 15\%$ ) we cannot directly see the edges imprint in the X–ray spectrum.125 Rather we are sensitive to the different slopes in between the edges leading to a precise determination of the depths of the single edges., Rather we are sensitive to the different slopes in between the edges leading to a precise determination of the depths of the single edges.126 We extract the data from an 80 pixel wide region. given," We extract the data from an 80 pixel wide region, given"127time of the observation.,time of the observation.128 The NIS light curve iu the 210 keV band indicates a rapid fiux increase by ~20% in the last 20 ksec exposure., The XIS light curve in the 2–10 keV band indicates a rapid flux increase by $\sim 20\%$ in the last 20 ksec exposure.129" The tine scale of this variability is ~LO! sec. indicating that tle cuission region is within ~30r, (ry.=cM is the eravitational radius) for a black holemass of M=—105ML. (sec section ??))."," The time scale of this variability is $\sim13010^4$ sec, indicating that the emission region is within $\sim 30131r_{\rm g}$ $r_{\rm g} \equiv \frac{GM}{c^2}$ is the gravitational radius) for a black holemass of $M = 10^{7.8} \solarmass$ (see section \ref{differ_SED}) )."132 By coutrast. the ποτ curve iu the 15LO keV. baud does not show evidence for significant variabilitv above the statistical errors iu the siue epoch.," By contrast, the light curve in the 15--40 keV band does not show evidence for significant variability above the statistical errors in the same epoch."133 Thus. there is a lint for decrease iu the harduess ratio between the PIN aud NIS. although its significance is mareial.," Thus, there is a hint for decrease in the hardness ratio between the PIN and XIS, although its significance is marginal."134 For the following spectral analysis. we separate the observation period iuto two. epoch 1 10? sec) and epoch 2 S10 LTLO? sec).," For the following spectral analysis, we separate the observation period into two, epoch 1 $\times 10^5$ sec) and epoch 2 $\times 10^5$ $\times 10^5$ sec)."135 We analyze the spectra of epoch 1 aud epoch 2 separately. by performing smiultaneous fit to those of the NIS-FIs in the 19 keV baud. the NIS-DI in the 15 keV baud. aud the IIND/PIN in the 1260 keV baud.," We analyze the spectra of epoch 1 and epoch 2 separately, by performing simultaneous fit to those of the XIS-FIs in the 1–9 keV band, the XIS-BI in the 1–8 keV band, and the HXD/PIN in the 12–60 keV band."136 Given the fact that the spectrum of 1€ 250.55 is subject to a heavy Galactic absorption. we do not utilize the NIS-BI data below 1 keV to avoid any possible uncertainties in the response.," Given the fact that the spectrum of 4C 50.55 is subject to a heavy Galactic absorption, we do not utilize the XIS-BI data below 1 keV to avoid any possible uncertainties in the response."137 We find that the effects of jle-up are significant above 9.0 keV for the NIS data., We find that the effects of pile-up are significant above 9.0 keV for the XIS data.138 Further. we discard the XIS data of the 1.71.9 keV baud jecause of calibration uncertainties associated with the iustruiental Si-lk edge.," Further, we discard the XIS data of the 1.7–1.9 keV band because of calibration uncertainties associated with the instrumental Si-K edge."139 Iu the sunultaucous fit. the flix iorinalizatious for the NIS-FIs and NIS-BI are set free. while that of the PIN is fixed at 1.18 relative to that of he NIS-FIs (Alaedaetal.2008).," In the simultaneous fit, the flux normalizations for the XIS-FIs and XIS-BI are set free, while that of the PIN is fixed at 1.18 relative to that of the XIS-FIs \citep{Mae08}."140. We firstly apply a single power law model Guocdified by he Galactic absorption fixed at 1.0«1022cm7) o the NIS aud PIN data in epochN= 1. covering the keV baud.," We firstly apply a single power law model (modified by the Galactic absorption fixed at $ = 1.0 \times 10^{22} \ {\rm cm^{-2}}$ ) to the XIS and PIN data in epoch 1, covering the 1--60 keV band."141 The fitis found to be far from acceptable qedof=13035/810) and results in a verv flat slope. Dx1.2.," The fitis found to be far from acceptable $\chi^2 / {\rm dof} = 13035/810$ ) and results in a very flat slope, $\Gamma \approx 1.2$."142 This is because the spectral shape below 10 keV ds πιο harder than that above 10 keV. Next we apply a cut-off power law model. in the form of E.P&exp(£f Leg). over which a narrow Catssian is added to represent an irou-I enuüssion liue at 6.1 keV. We obtain DzmO83 and EaocLl keV with \?/dof=1616/808.," This is because the spectral shape below 10 keV is much harder than that above 10 keV. Next we apply a cut-off power law model, in the form of $E^{-{\rm \Gamma}} \times {\rm exp}(-E/E_{\rm cut})$ , over which a narrow Gaussian is added to represent an iron-K emission line at 6.4 keV. We obtain ${\rm \Gamma} \approx 0.83$ and $E_{\rm cut} \approx14314$ keV with $\chi^2 /{\rm dof} = 4616/808$ ."144 There remain stroug absorption features around 1 keV sugecsting intrinsic absorption. however.," There remain strong absorption features around 1 keV suggesting intrinsic absorption, however."145 When we add another absorption at the source redshift (+=0.02). the fit becomes much better. vielding P—1.3945:0.02. Loy=1542 keV. and Ny=(0.67£0.01)&10°?emn7 with κfdof=997/807.," When we add another absorption at the source redshift $z=0.02$ ), the fit becomes much better, yielding $\Gamma146=1.39\pm0.02$, $E_{\rm cut} = 45\pm2$ keV, and $N_{\rm H} =147(0.67\pm0.01) \times 10^{22} \ {\rm cm^{-2}}$ with $\chi^2/{\rm dof} =148997/807$."149 Finally. similar to the analysis done bv Molina.etal. (2007).. we consider double absorber model with two differeut column deusities. aadNg. whose covering fraction is f and (1—f). respectively.," Finally, similar to the analysis done by \citet{Mol07}, we consider double absorber model with two different column densities, and, whose covering fraction is $f$ and $(1-f)$, respectively."150" The fit is sienificautly improved (\7/dof= 910/805) with Po—L6L'e"" E=ll! keV. ΓΕ. ο(73«1072cin7. aud f=O19."," The fit is significantly improved $\chi^2/{\rm dof} =151940/805$ ) with $\Gamma = 1.61^{+0.03}_{-0/05}$, $E_{\rm cut} = 140^{+46}_{-20}$ keV, $ \approx 8.3\times 10^{22} \ {\rm cm^{-2}}$, $ \approx 0.73 \times 10^{22} \ {\rm cm^{-2}}$, and $f= 0.19$."152" The photon iudex becomes a reasonable value for Αννα,", The photon index becomes a reasonable value for AGNs.153" Note that our ""double absorber” model has three free parameters for the absorber. while the ""double partial covering"" iodel in the NSPEC terminology) adopted by Molinactal.(2007) has four."," Note that our “double absorber” model has three free parameters for the absorber, while the “double partial covering” model in the XSPEC terminology) adopted by \citet{Mol07} has four."154 Since the latter model docs not eive a significant improvement for our data (A4?z 1) we adopt the former as a base contiuuuni model for the following analysis.," Since the latter model does not give a significant improvement for our data $\Delta \chi^2 \approx 1$ ), we adopt the former as a base continuum model for the following analysis."155 The high qualitySneak spectra are quite useful to constrain the reflection component. which is iudicated by the presence of the iron-I& emission line.," The high quality spectra are quite useful to constrain the reflection component, which is indicated by the presence of the iron-K emission line."156" Thus. we include it bv ufiliziug our modified version of ""pexiiv reflection code CMagdziarz&Zdzimski1995) that assumes a cutoff power law continui aud contains a selt-consisteut fluorescence dron-E line calculated according to the same aleorithimi as described im Zyvekietal.(1999)."," Thus, we include it by utilizing our modified version of “pexriv” reflection code \citep{Mag95} that assumes a cutoff power law continuum and contains a self-consistent fluorescence iron-K line calculated according to the same algorithm as described in \citet{Zyc99}."157. The additional free paraimcter is the relative reflection streugth. πι O/2a). whilewe fix theionization parameter. temperature. and inclination angle at 0.10° Ik. and 35 degrees (Molimaetal. 2007)... respectively.," The additional free parameter is the relative reflection strength, $R (\equiv \Omega / 2\pi)$ , whilewe fix theionization parameter, temperature, and inclination angle at 0,$10^5$ K, and 35 degrees \citep{Mol07}, , respectively."158 The solar Fe abuudauce by Aucers&Crevesse(1989). (Fe/TI = LGS&10 7) is asstuned., The solar Fe abundance by \citet{and89} (Fe/H = $4.68\times10^{-5}$ ) is assumed.159 Considering that the reflection, Considering that the reflection160parity asvmmet(iry and compare them with the CAIB kinematic dipole.,parity asymmetry and compare them with the CMB kinematic dipole.161 In Section 4. sunmiarize our investigation.," In Section 4, we summarize our investigation."162" The temperature fIuctuations of CMD anisotropy. can be conveniently. decomposed as follows: where @,, are the coefficients of decomposition: ej,=μηexp(opn,). with on, as the phase."," The temperature fluctuations of CMB anisotropy, can be conveniently decomposed as follows: where $a_{lm}$ are the coefficients of decomposition: $a_{lm} = |a_{lm}| \exp(i \phi_{lm})$ , with $\phi_{lm}$ as the phase."163 Under (he assumption of total Gaussian randonmness. as predicted by the large class of inflationary models. the amplitudes αμ are distributed according to Ravleigh’s probability distribution function aud the phases of αμ are supposed to be evenly distributed in the range [0.27] (Bardeenetαἱ.1986).," Under the assumption of total Gaussian randomness, as predicted by the large class of inflationary models, the amplitudes $|a_{lm}|$ are distributed according to Rayleigh's probability distribution function and the phases of $a_{lm}$ are supposed to be evenly distributed in the range $[0,2 \pi]$ \citep{bardeen1986}."164". For anv signals Z(n) defined on the sphere. one can extract svmmetric (=n)) and antisvimmetric (AT(n)=—AT (—n)) components. where and DP(/)=cos?(ay DL(/)=sin?(#). τα)=(-1)Y5,(—n)."," For any signals $T(\hat{\mathbf n})$ defined on the sphere, one can extract symmetric $ \Delta T^+(\hat{\mathbf n})=\Delta T^+(-\hat{\mathbf n} $ )) and antisymmetric $ \Delta T^-(\hat{\mathbf n}) =- \Delta T^-(-\hat{\mathbf n})$ ) components, where and $\Gamma^{+}(l)\equiv\cos^2(\frac{\pi l}{2})$, $\Gamma^{-}(l)\equiv\sin^2(\frac{\pi l}{2})$, $Y_{lm}(\hat{\mathbf n})=(-1)^l\,Y_{lm}(-\hat{\mathbf n})$."165 Naive expectation. where the concordant ACDAL cosmological model with initial statistically isotropic and Gaussian acdiabatie perturbations is assumed. is the absence of any features distinct between even and odd multipoles.," Naive expectation, where the concordant $\Lambda$ CDM cosmological model with initial statistically isotropic and Gaussian adiabatic perturbations is assumed, is the absence of any features distinct between even and odd multipoles."166 ILowever. in reality this statement needs more accurate clarification.," However, in reality this statement needs more accurate clarification."167 In particular. for the curvature perturbations bevond (he present horizon the power spectrum is given by P(E)xfhors where ηςo90.96 is the spectral index of the primordial density perturbations (IXomatsuetal. 2011)..," In particular, for the curvature perturbations beyond the present horizon the power spectrum is given by $P(k)\propto k^{-4+n_s}$, where $n_s\simeq 0.96$ is the spectral index of the primordial density perturbations \citep{komatsu2011}. ."168" Thus. the variance of the metric perturbations 9?[1pPU.MxKstin)1 Has very weak power-law (0,<» 1)ς or logarithmic. divergence. (0,c 1).x ilaE Αμάν—0."," Thus, the variance of the metric perturbations $\sigma^2\sim\int k^2P(k)dk\propto k_{\min}^{n_s-1}$ has very weak power-law $n_s<1$ ) or logarithmic divergence $n_s\simeq 1$ ), if $k_{\min}\rightarrow 0$."169 MSince the low multipole range of the CAMB temperature anisotropy is determined by the and integrated Sachs-Wolle effects. (hese peculiarity of the power spectrum of metric perturbations arecrucial for the two-point correlation function:," Since the low multipole range of the CMB temperature anisotropy is determined by the ordinaryand integrated Sachs-Wolfe effects, these peculiarity of the power spectrum of metric perturbations arecrucial for the two-point correlation function:"170the most recent study which claims a factor of four euliaucemeut on the nearsidle when compared to the far (LeFeuvreaud.Wieczorek2005).. but are in good agreement with the work cdoue by Wiesel(1971) ancl BanclerinannaucSinger(1973).,"the most recent study which claims a factor of four enhancement on the nearside when compared to the far \citep{fev05}, but are in good agreement with the work done by \citet{wiesel71} and \citet{band73}."171. We see mild evidence for Baucermannu ancl Singer's (1973) assertion of the Earth acting as a shield for «258. and little effect outside this distance., We see mild evidence for Bandermann and Singer's (1973) assertion of the Earth acting as a shield for $a < 25 R_\oplus$ and little effect outside this distance.172 Thus. Bandermaun and Singers estimate of no measurable nearside/farside asymunetry in the current cratering rate is correct.," Thus, Bandermann and Singer's estimate of no measurable nearside/farside asymmetry in the current cratering rate is correct."173 Iu the bulk of our simulations we have used the approximation that the Moou’s orbit is iu the ecliptic plane., In the bulk of our simulations we have used the approximation that the Moon's orbit is in the ecliptic plane.174 Since we show that the latituclinal dependence of Iunar cratering is weak (~ reduction within of the poles relative to a baud centered on the Moou's equator). we do uot expect the inclusion of the moon's orbital inclination to alter our results siguificautly. although we expect the polar asymmetry to monotonically decrease with increasing orbital inclination.," Since we show that the latitudinal dependence of lunar cratering is weak $\sim$ reduction within $^{\circ}$ of the poles relative to a band centered on the Moon's equator), we do not expect the inclusion of the moon's orbital inclination to alter our results significantly, although we expect the polar asymmetry to monotonically decrease with increasing orbital inclination."175 We have confirmed this by computing the CALAACA aud polar asyuunetry ratios [or a less-exteusive set of simulations in which the lunar orbital inclination is initially set to its current value of 5.15° aud we use the sub-Earth point at the time ofimpact to compute lunocentric latitudes ancl longitudes., We have confirmed this by computing the GMAACA and polar asymmetry ratios for a less-extensive set of simulations in which the lunar orbital inclination is initially set to its current value of $^{\circ}$ and we use the sub-Earth point at the time of impact to compute lunocentric latitudes and longitudes.176 We find a slight reduction of GNLAAC'A to 1.212:0.02 from (1.292 0.01) auc a crater density. within of the pole that is statistically the same as the inclination case (0.911+0.009 instead of 0.912+0.00 1)., We find a slight reduction of GMAACA to $1.24\pm0.02$ from $1.29\pm0.01$ ) and a crater density within $^{\circ}$ of the pole that is statistically the same as the inclination case $0.914 \pm 0.009$ instead of $0.912 \pm 0.004$ ).177 We have used the debiased NEO moclel of Bottkeefa£.(2002) to examine the bombarcimer ol the Earth-Moon system in terms of various impact aud crater asvunmetries., We have used the debiased NEO model of \citet{bottke02} to examine the bombardment of the Earth-Moon system in terms of various impact and crater asymmetries.178 For Earth arrivals we finda «1% variation in the ratio between the areal densities within of the poles and within a baud centered on the equator., For Earth arrivals we find a $< 1$ variation in the ratio between the areal densities within of the poles and within a band centered on the equator.179 The local time distribution of terrestrial impacts from NEOs is enhanced duriug the AM hours., The local time distribution of terrestrial impacts from NEOs is enhanced during the AM hours.180 While this fall-time distribution corresponds well to recent radar data. it is in disagreement with the choudritic meteorite data aud their derived. pre-atimospheric orbital distributions.," While this fall-time distribution corresponds well to recent radar data, it is in disagreement with the chondritic meteorite data and their derived pre-atmospheric orbital distributions."181 This discrepaney thus reinforces the conclusion of Morbidelli(1998) that the large amount of clecimeter-scale material being ejected from the main asteroid be outo Earthi-crossiug orbits must be collisionally depleted before much of it cau evolve to orbits with qc1.5 AU., This discrepancy thus reinforces the conclusion of \citet{morglad98} that the large amount of decimeter-scale material being ejected from the main asteroid belt onto Earth-crossing orbits must be collisionally depleted before much of it can evolve to orbits with $a < 1.5$ AU.182 A significant result is that we find the average impact speed onto the Moon to be km/s. with a non-negligible hieher-speed tail (Fig. 11)).," A significant result is that we find the average impact speed onto the Moon to be $\bar{v}_{imp} = 20$ km/s, with a non-negligible higher-speed tail (Fig. \ref{fig:velhist}) )."183 This combined with quantification of the nou-uniform surface cratering has iaplicatious for both tracing crater fields back to the size distribution of the impactors an he absolute (or relative) dating of cratered surfaces., This combined with quantification of the non-uniform surface cratering has implications for both tracing crater fields back to the size distribution of the impactors and the absolute (or relative) dating of cratered surfaces.184 First. the higher impact speeds we find mean that lunar impact craters (at least in tlie post-inare era when we believe the NEO orbital distribution we are using is valid) have been produced by smaller inpactors than previously caleulated.," First, the higher impact speeds we find mean that lunar impact craters (at least in the post-mare era when we believe the NEO orbital distribution we are using is valid) have been produced by smaller impactors than previously calculated."185 This τςighly, This roughly186simulator to replicate conditions present at the telescope.,simulator to replicate conditions present at the telescope.187"constraints on the spatial distribution of the mmolecules, we have used data obtained with the VLA that will be published separately (Menten et al.","constraints on the spatial distribution of the molecules, we have used data obtained with the VLA that will be published separately (Menten et al."188 in prep.)., in prep.).189 The hfs splitting in the inversion lines in velocity units is much wider than for the rotation line., The hfs splitting in the inversion lines in velocity units is much wider than for the rotation line.190" However, because both the (1,1) and (2,2) lines are very optically thin, as indicated by the spectra and supported by our modeling (see Sect.3.2)),"," However, because both the (1,1) and (2,2) lines are very optically thin, as indicated by the spectra and supported by our modeling (see Sect.\ref{modeling}) ),"191 any contribution of the hfs components will be factors of several weaker than the main hfs component and neglected in the modeling., any contribution of the hfs components will be factors of several weaker than the main hfs component and neglected in the modeling.192" The intensities of the spectra produced from the VLA images, which were restored with a circular beam of 3.7 FWHM, used in that section, are consistent with the published 100 m telescope values."," The intensities of the spectra produced from the VLA images, which were restored with a circular beam of $3\as7$ FWHM, used in that section, are consistent with the published 100 m telescope values."193" InFigs. 1,,"," InFigs. \ref{nh3plot}, ,"194" 2 and Table 1 we present the results of our HIFI observations of both the aand the H2O ortho ground state lines, together with the results of our mmodeling."," \ref{vlaplot} and Table \ref{lineresults} we present the results of our HIFI observations of both the and the $_2$ O ortho ground state lines, together with the results of our modeling."195" All our observed positions agree to within 2"" with the stars’ 2MASS positions, which themselves have an absolute accuracy of better than 0/1 (Cutrietal.2003)."," All our observed positions agree to within $2''$ with the stars' 2MASS positions, which themselves have an absolute accuracy of better than $0\as1$ \citep{Cutri2003}."196. The determination of the LSR ranges is somewhat subjective and the upper and lower velocities are uncertain by ~ a few ffor weaker lines., The determination of the LSR ranges is somewhat subjective and the upper and lower velocities are uncertain by $\sim$ a few for weaker lines.197" For all entries, the formal error in fTmpdv is smaller than 0.1 Ks~!.. For VY CMa we used the higher of the literature mass-loss rate values scaled to the recently measured trigonometric parallax distance, 1100 pc (Choiet 2008)."," For all entries, the formal error in $\int T_{\rm MB}d$$\varv$ is smaller than 0.1 K. For VY CMa we used the higher of the literature mass-loss rate values scaled to the recently measured trigonometric parallax distance, 1100 pc \citep{Choi2008}."198". For both IK Tau and IRC+10420 the lower and the higher values of Χνη, are implied by the cm lines and the submm line, respectively."," For both IK Tau and IRC+10420 the lower and the higher values of $X_{{\rm NH}_3}$ are implied by the cm lines and the submm line, respectively."199 Inspecting Fig., Inspecting Fig.200" | and the table, it is striking to see that the luminosity (integrated intensity) in the aand H5O ground-state lines is of comparable magnitude for all of our objects."," \ref{nh3plot} and the table, it is striking to see that the luminosity (integrated intensity) in the and $_2$ O ground-state lines is of comparable magnitude for all of our objects."201" The NH3//H5O line ratios are «0.50,0.30,0.28, and 0.42 for IK Tau, VY CMa, OH 26.5+0.6, and IRC+10420, respectively."," The $_2$ O line ratios are $\approx 0.50, 0.30, 0.28,$ and 0.42 for IK Tau, VY CMa, OH 26.5+0.6, and IRC+10420, respectively."202" One has to keep in mind that H,O is a major molecular constituent of our CSEs, while even the presence of observable eemission is completely unexplained!"," One has to keep in mind that $_2$ O is a major molecular constituent of our CSEs, while even the presence of observable emission is completely unexplained!"203" Another remarkable result is that the velocity ranges covered by the two lines are almost identical, which suggests that the bulk of the material producing the emission for both is similar."," Another remarkable result is that the velocity ranges covered by the two lines are almost identical, which suggests that the bulk of the material producing the emission for both is similar."204" Moreover, for all our targets, both lines FWZP values are lower, but comparable to twice the terminal velocity, implying that both molecules are present in the outer layers of the envelope, where the material has almost been fully accelerated."," Moreover, for all our targets, both lines' FWZP values are lower, but comparable to twice the terminal velocity, implying that both molecules are present in the outer layers of the envelope, where the material has almost been fully accelerated."205" Furthermore, we point out the clear self absorption in the blue wing of the H2O lines toward TTau and CCMa, which proves that the line emitting region covers the envelope."," Furthermore, we point out the clear self absorption in the blue wing of the $_2$ O lines toward Tau and CMa, which proves that the line emitting region covers the envelope."206 Whether this is also true for the line is a priori not clear., Whether this is also true for the line is a priori not clear.207 The eemission of the sources has been modeled with the Monte Carlo radiative transfer code RATRAN developed and described by Hogerheijde&vanderTak(2000)., The emission of the sources has been modeled with the Monte Carlo radiative transfer code RATRAN developed and described by \citet{hogerheijde2000}.208". For NH3,, RATRAN uses collision rates calculated by Danbyetal. (1988)."," For , RATRAN uses collision rates calculated by \citet{Danby1988}."209". Power laws for the density and temperature were usedto describe the physical structure of the envelope, using as input published values for the mass loss rate and expansion velocity (see Table 1))."," Power laws for the density and temperature were usedto describe the physical structure of the envelope, using as input published values for the mass loss rate and expansion velocity (see Table \ref{lineresults}) )."210" For IRC+10420, envelope parameters from Dinh-V.-Trungetal.(2009) were used and for IK Tau, VY CMa, and OH26.5+0.6 we refer to the modeling of Decin (2010),, Decinetal. (2006),, andJusttanontetal. (2006),, respectively."," For IRC+10420, envelope parameters from \citet{trung2009} were used and for IK Tau, VY CMa, and OH26.5+0.6 we refer to the modeling of \citet{Decin2010}, , \citet{Decin2006}, , and\citet{Justtanont2006}, , respectively."211" The VLA data yield an extent of «4"" for the", The VLA data yield an extent of $\approx 4''$ for the212primary continuum.,primary continuum.213 The final fit is good (ντ=249/208 cLo..):, The final fit is good $\chi^2=249/208$ d.o.f.):214 the continuum parameters are shown in ‘Table 1.. while a detailed discussion on the emission line spectrum is elerredto Sect. 4..," the continuum parameters are shown in Table \ref{continuum}, while a detailed discussion on the emission line spectrum is deferredto Sect. \ref{discussion}. ."215 Fig., Fig.216 2(a) and 2(b) show the spectrum uxd the adopted. best fit mocel., \ref{pnspectrum} and \ref{emodel} show the spectrum and the adopted best fit model.217" The powerlaw index of the primary continuum. and the local absorbing column density are well constrained. the (wo values being. L7μα+0.01 and L36speULES($5,«107!2 7."," The powerlaw index of the primary continuum and the local absorbing column density are well constrained, the two values being $1.77\pm0.01$ and $1.36^{+0.03}_{-0.04}\times10^{24}$ $^{-2}$."218 An interesting result from the fit is an hint [or iron underabundance. the best fit⋅ value being. 0.82.κOde5122 (withHre-specttothesolarabundancesmeasuredby 2).," An interesting result from the fit is an hint for iron underabundance, the best fit value being $0.82^{+0.10}_{-0.08}$ \citep[with respect to the solar abundances measured by][]{ag89}."219 However. this result must be taken with caution. because this parameter Is actually calculated on the basis of the iron edge depth of the Compton reflection continuum. which depends tightly also on the primary powerlaw index and the inclination angle.," However, this result must be taken with caution, because this parameter is actually calculated on the basis of the iron edge depth of the Compton reflection continuum, which depends tightly also on the primary powerlaw index and the inclination angle."220 The Compton reflection component and the strong iron he ine are clear signatures of reflection from a Compton-thick. fairly neutral. material.," The Compton reflection component and the strong iron $\alpha$ line are clear signatures of reflection from a Compton-thick, fairly neutral, material."221 It is then natural to assume that the absorber and the rellector are one and the same material. he torus. as usually found in other Compton-thick Sevfert ealaxies.," It is then natural to assume that the absorber and the reflector are one and the same material, the torus, as usually found in other Compton-thick Seyfert galaxies."222 Ht. is interesting to note that we also detected a ine consistent with emission from Si less ionised than (sce Table 23). thus likely being another product of reflection rom the torus.," It is interesting to note that we also detected a line consistent with emission from Si less ionised than (see Table \ref{toruslines}) ), thus likely being another product of reflection from the torus."223 This feature was also found in the spectrum (?).., This feature was also found in the spectrum \citep{sako00b}.224 The 2-10 keV tus of the Compton reflection component is 3.410.* erg 27s + completely consistent with the 33.10 ας em7E + measured by the observation. performed seven months before.," The 2-10 keV flux of the Compton reflection component is $3.4\times10^{-12}$ erg $^{-2}$ $^{-1}$, completely consistent with the $3.3\times10^{-12}$ erg $^{-2}$ $^{-1}$ measured by the observation, performed seven months before."225 To check for variability in past X-ray. observations. we can compare the 2-10 keV flux reported in ‘Table 1.. which by far is dominated by the rellection component (see Fig. 2(b))).," To check for variability in past X-ray observations, we can compare the 2-10 keV flux reported in Table \ref{continuum}, , which by far is dominated by the reflection component (see Fig. \ref{emodel}) )."226 The 1997 BeppoSAX observation measured a Πας of 6.5«10.72 erg 7s + (9?) slightly larger than the EPIC pn one. the dillerence probably Iving in the brighter intrinsic continuum of the source at that time (see below).," The 1997 BeppoSAX observation measured a flux of $6.5\times10^{-12}$ erg $^{-2}$ $^{-1}$ \citep{cappi99}, , slightly larger than the EPIC pn one, the difference probably lying in the brighter intrinsic continuum of the source at that time (see below)."227 Indeed. ο performed a detailed comparison between their dataset and the and observations. concluding that there was evidence of a variable component above the iron line energy. and a constant one below.," Indeed, \citet{cappi99} performed a detailed comparison between their dataset and the and observations, concluding that there was evidence of a variable component above the iron line energy and a constant one below."228 This is also supported. by some indications that the intrinsic. unabsorbed luminosity (O.1-150 keV) of AIrk 3 changed between the ancl observations and with respect to the BeppoSAX one. the values being 1.10. os)107 and L3.107 ere sd...," This is also supported by some indications that the intrinsic, unabsorbed luminosity (0.1-150 keV) of Mrk 3 changed between the and observations and with respect to the BeppoSAX one, the values being $1.1\times10^{44}$, $0.8\times10^{44}$ and $1.3\times10^{44}$ erg $^{-1}$,."229 However. it should be stressed that this luminosity is driven by the normalization of the stronglv absorbed. powerlaw. so the measures with and. in particular. are much less reliable than the one obtained with the broadband spectrum of BeppoSAX.," However, it should be stressed that this luminosity is driven by the normalization of the strongly absorbed powerlaw, so the measures with and, in particular, are much less reliable than the one obtained with the broadband spectrum of BeppoSAX."230 Finally. an estimate of the angle 7 can be mace on the basis of the amount of Compton rellection with respect to the incidentnuclear continuum which is quite large (2 7).. thus requiringan inclination not very high toavoid of the torus.," Finally, an estimate of the angle $i$ can be made on the basis of the amount of Compton reflection with respect to the incidentnuclear continuum which is quite large \citep[$R\simeq1$:][]{cappi99}, , thus requiringan inclination not very high toavoid self-obscuration of the torus."231whether ieionization Was occurmiue from the fforest.,whether reionization was occurring from the forest.232 Appendix À outlines how to measure the hardness of the ionizing background with aabsorption. and Appendix 3. derives formulae for the sienificance with which aabsorptiou can be detected iu cross correlation with the coeval Lya forest.," Appendix \ref{sec:hardness} outlines how to measure the hardness of the ionizing background with absorption, and Appendix \ref{ap:SNRcc} derives formulae for the significance with which absorption can be detected in cross correlation with the coeval $\alpha$ forest."233" This paper assumes a flat ACDAL cogsnologv with h—QT. Q,=0.016. O,,=(28. σε=0.39. n,—I. and Vip.=0.21. consistent with recent measurements (I&omatsuetal. 2009)."," This paper assumes a flat $\Lambda$ CDM cosmology with $h =0.7$, $\Omega_b = 0.046$, $\Omega_m = 0.28$, $\sigma_8 = 0.82$ $n_s = 1$ , and $Y_{\rm He} = 0.24$, consistent with recent measurements \citep{komatsu08}. ."234". ILowever. the simulation used to calculate the Ίσα forest spectra at 2<1.5. the D5 simulation iu Springel&Ieruquist(2003).. asses a slightly different cosmology. with the most notable differences being Q,=001 and o4=0.9."," However, the simulation used to calculate the $\alpha$ forest spectra at $z < 1.5$, the D5 simulation in \citet{springel03}, assumes a slightly different cosmology, with the most notable differences being $\Omega_b = 0.04$ and $\sigma_8 = 0.9$."235 The photoionization and recombination rates used in this study are from Hui&Cuecin(1997)., The photoionization and recombination rates used in this study are from \citet{hui97}.236. Iu photoionization equilibriun. the ffraction is determined by the relation Photoionization οὐαυτα is a good approximation because the timescaleto reach equilibrium vr at τοGh ego Faucho-Cüguereetal. 2008a)) was much shorter than all other relevant timescales.," In photoionization equilibrium, the fraction is determined by the relation Photoionization equilibrium is a good approximation because the timescaleto reach equilibrium $\Gamma_{\rm HeI}^{-1} \sim \Gamma_{\rm HI}^{-1} \approx 3 \times 10^4$ yr at $z\sim 3$; e.g., \citealt{faucher08}) ) was much shorter than all other relevant timescales."237 Here. ay. ny. avy. aud Dy=πριΗΕE)ox(E)JCE) are respectively the Case A recombination the mumuber density. the ionization fraction. aud the photoionization rate for ion X (or subscript e for electrons). and aud oy are the ionizatiou potential aud the photoionizatiouEw cross section.," Here, $\alpha_{X}$, $n_{X}$, $x_X$ , and $\Gamma_{X} \equiv \int_{E^{\rm ion}_X}^\infty (dE/E) \, \sigma_X(E) \, J(E)$ are respectively the Case A recombination the number density, the ionization fraction, and the photoionization rate for ion $X$ (or subscript $e$ for electrons), and $E^{\rm ion}_X$ and $\sigma_X$ are the ionization potential and the photoionization cross section."238 Lastly. (E) is the incident specific intensity inteerated over solid angle.," Lastly, $J(E)$ is the incident specific intensity integrated over solid angle."239" We will sometime:μα imupreciselv write μοι=d for a gas parcel even thoug[um a sinall fraction (~10.7 A,) of the Ποια is in cnr."," We will sometimes imprecisely write $x_{\rm HeII} = 1$ for a gas parcel even though a small fraction $\sim 10^{-5}\, \Delta_b$ ) of the helium is in $x_{\rm HeI}$."240 If the cleusity aud Py are known (sinee n. ds effectively known iu the IGM after hydrogen reionizatiou up to denusitv). a measurement of πω cau be used to determine wy (equation 1)).," If the density and $\Gamma_{\rm HeI}$ are known (since $n_e$ is effectively known in the IGM after hydrogen reionization up to density), a measurement of $x_{\rm HeI}$ can be used to determine $x_{\rm HeII}$ (equation \ref{eqn:xHeI}) )."241 Fortunately. coeval LLxauan-series absorption provides an estunate for the density of a aabsorber.," Fortunately, coeval Lyman-series absorption provides an estimate for the density of a absorber."242" In addition. the pphotoionization rate (like this for the 1)) is expected to have been esseutiallv spatially incependeut owing to the long mean free path of t-ionizing photous (1,4) and the large umber of sources iu a volume of ~ I"," In addition, the photoionization rate (like this for the ) is expected to have been essentially spatially independent owing to the long mean free path of -ionizing photons $l_{\rm HeI}$ ) and the large number of sources in a volume of $\sim l_{\rm HeI}^3$."243t is ostimated that μοιo>7j9300 cAIpe at 2~3 (MadauTE 20093.," It is estimated that $l_{\rm HeI} > l_{\rm HI} \sim 300~$ cMpc at $z \sim 3$ \citep{madau99, faucher08, prochaska09}."244. Therefore. Pupp was effectively just a single ΠΡΟ: in all of the IGAL," Therefore, $\Gamma_{\rm HeI}$ was effectively just a single number in all of the IGM."245" This paper focuses primarily on the longesSF.+ wavelength allowed erouud-state transition forr.. theIP"". 581 ttrausition."," This paper focuses primarily on the longest wavelength allowed ground-state transition for, the, $584~$ transition."246 Its oscillator strength is 3.8 times larger than the next strougest transition ofLr. the!1DU. 537 ttrausition.," Its oscillator strength is $3.8$ times larger than the next strongest transition of, the, $537~$ transition."247 Hala aabsorptionisobsercalbleinasuitablequasarabsorptionspectrumbetwe aud 3OL(L|2030)A. where the LLwa forest begins.," $584~$ absorption is observable in a suitable quasar absorption spectrum between the wavelengths of $584 \, (1+z_{\rm QSO})~$ and $304\, (1+z_{\rm QSO})~$ , where the $\alpha$ forest begins."248" This wavelength rauge correspouds to aabsorption from 0.5soso<2Xtoso- To redshift across the 581 ttransition. a photon experienced the optical depth from eas in the IIubble flow of (hat is termed the CCumau-Petersou optical depth: Cama&Trippetal.1990)) where is the Live Cunu-Poeterson optical depth. Aj is the gas density in units of the cosuuc mean. Z7, is the temperature im units of 10! K. aud the 0.7 exponeut owes to the temperature dependence of the recombination cocfiicicut (c£."," This wavelength range corresponds to absorption from $0.5 \, z_{\rm QSO} <z < z_{\rm QSO}$ To redshift across the $584~$ transition, a photon experienced the optical depth from gas in the Hubble flow of (what is termed the Gunn-Peterson optical depth; \citealt{gunn65,tripp90}) ) where is the $\alpha$ Gunn-Peterson optical depth, $\Delta_b$ is the gas density in units of the cosmic mean, $T_4$ is the temperature in units of $10^4~$ K, and the $-0.7$ exponent owes to the temperature dependence of the recombination coefficient (cf."249 equation 1))., equation \ref{eqn:xHeI}) ).250 Overdensities of a few aud greater at 2~23 hac decoupled from the ΠΠολο flow aud were collapsing or had collapsed., Overdensities of a few and greater at $z \sim 3$ had decoupled from the Hubble flow and were collapsing or had collapsed.251 Tn these regious. the Camu-Petersou optical depth no longer describes the absorption.," In these regions, the Gunn-Peterson optical depth no longer describes the absorption."252 Iustead. such resions appear as distinct absorption lues with widths of 10s of lau 1. and tbeir optical depth iu SalaHisto where LM is the LLwe optical depth. and we have approximated the line profile as a tophat with velocity widthAcy (aud lxἈυμιἈυμα< 2. with 2 being the liitof pure thermal broadening).," Instead, such regions appear as distinct absorption lines with widths of $10$ s of km $^{-1}$, and their optical depth in $584~$is where $\tau_{\rm HI, 1216}^{N_{\rm HI}}$ is the $\alpha$ optical depth, and we have approximated the line profile as a tophat with velocity width$\Delta v_X$ (and $1 \leq {{\Delta v}_{\rm HI}}/{{\Delta v}_{\rm HeI}} \leq 2$ , with $2$ being the limitof pure thermal broadening)."253 Systems with ΑΠ>Lothem? are common in the Lya forest.," Systems with $N_{\rm HI} > 10^{14} \, {\rm cm}^{-2}$ are common in the $\alpha$ forest."254 Each sightline intersects 100 such systems between 2=3 aud:= L(esg.. 1993)).," Each sightline intersects $\sim 100$ such systems between $z=3$ and $z=4$ (e.g., \citealt{press93}))."255 The hala oopticaldepthdepondsonthecalucofl yg wa addition to the feld of mterest. yp.," The $584$ optical depth depends on the value of $\Gamma_{\rm HeI}$ in addition to the field of interest, $x_{\rm HeII}$ ."256 The calculatious in this paper asstune Dg= Py.and Py is chosen to match nieasuremients of the LLva meu transmission (requiring Typ~100 8 k," The calculations in this paper assume $\Gamma_{\rm HeI} = \Gamma_{\rm HI}$ ,and $\Gamma_{\rm HI}$ is chosen to match measurements of the $\alpha$ mean transmission (requiring $\Gamma_{\rm HI} \sim 10^{-12}~$ $^{-1}$ ;"257 The calculatious in this paper asstune Dg= Py.and Py is chosen to match nieasuremients of the LLva meu transmission (requiring Typ~100 8 kk," The calculations in this paper assume $\Gamma_{\rm HeI} = \Gamma_{\rm HI}$ ,and $\Gamma_{\rm HI}$ is chosen to match measurements of the $\alpha$ mean transmission (requiring $\Gamma_{\rm HI} \sim 10^{-12}~$ $^{-1}$ ;"258Emission at 7 keV in the source frame would require the iron to be hiehly ionized.,Emission at 7 keV in the source frame would require the iron to be highly ionized.259 The photon indices for the soft and hard. power laws are 3.47 and 0.99. respectively.," The photon indices for the soft and hard power laws are 3.47 and 0.99, respectively."260 Figure 6 shows the observed spectrum of Z11593-0112 together with the model components., Figure \ref{fig:z11598} shows the observed spectrum of Z11598-0112 together with the model components.261 The significance of the line cannot be tested using the F-test because the test is only valid for Gaussian statistics., The significance of the line cannot be tested using the F-test because the test is only valid for Gaussian statistics.262 Therefore. simmlated “lake” spectra. constructed in NSPEC.. were used (o determine the likelihood that the emission line seen in Z11598-0112 is real.," Therefore, simulated “fake” spectra, constructed in XSPEC, were used to determine the likelihood that the emission line seen in Z11598-0112 is real."263 A set of 500 fake spectra were created using the FANETT command in XSPEC., A set of 500 fake spectra were created using the FAKEIT command in XSPEC.264 The task uses the response matrices associated with the real spectra aud the best-fit continuum moclel to create artificial source ancl background spectra., The task uses the response matrices associated with the real spectra and the best-fit continuum model to create artificial source and background spectra.265 The simulated source and backerounedl spectra were also binned to at least 3 counts per bin., The simulated source and background spectra were also binned to at least 3 counts per bin.266 As a sanity check. we modeled the simulated spectra.," As a sanity check, we modeled the simulated spectra."267 The distribution of the photon indices in the 500 simulated spectra were consistent with the distribution of the photon indices for the observation of Z11598-0112. modeled with 3 counts per bin.," The distribution of the photon indices in the 500 simulated spectra were consistent with the distribution of the photon indices for the observation of Z11598-0112, modeled with 3 counts per bin."268 Therefore. we are confident that binning the data to an arbitrary siall nunber of counts per bin did not introduce anv biases.," Therefore, we are confident that binning the data to an arbitrary small number of counts per bin did not introduce any biases."269 We lound (hat only 3 ol the 500 spectra showed a flux at the energy of the line exceeding (he measured line fIux minus its error bar (a conservative measure of the line flux)., We found that only 3 of the 500 spectra showed a flux at the energy of the line exceeding the measured line flux minus its error bar (a conservative measure of the line flux).270 Thus. the line is significant ad above the level.," Thus, the line is significant at above the level."271 It is also important to note that Z11598-0112 is considered to be a Narrow-Line Sevfert 1 (NLS1) galaxy whose LL? line width (FWIIM) is 770 km |. based on the data presented in Veilleuxetal.(1999a).," It is also important to note that Z11598-0112 is considered to be a Narrow-Line Seyfert 1 (NLS1) galaxy whose ${\beta}$ line width (FWHM) is 770 km $^{-1}$, based on the data presented in \citet{vei99a}."272. NLSIs tend to have steeper soft X-ray spectra than normal Sevlert 1 ealaxies. as shown independentlv by observations analyzed bv Leiehlv(1999) and Vaughanetal.(1999).," NLS1s tend to have steeper soft X-ray spectra than normal Seyfert 1 galaxies, as shown independently by observations analyzed by \citet{leighly} and \citet{vaughan}."273. Of the 24 NLSIs studied by Leighly(1999). and. (1999).. have soft [lux in excess of the power law mocdel that fits the individual spectra at high energies: the excess flux dominates the spectra at energies < 1.5 keV. The nominal power law photon indices of NLSIs over the 0.610 keV energv baud span the range of 1.62.5. lareer than normal Sevlert 1 galaxies (Vaughanetal.1999).," Of the 24 NLS1's studied by \citet{leighly} and \citet{vaughan}, have soft flux in excess of the power law model that fits the individual spectra at high energies; the excess flux dominates the spectra at energies $\lesssim$ 1.5 keV. The nominal power law photon indices of NLS1s over the 0.6–10 keV energy band span the range of 1.6–2.5, larger than normal Seyfert 1 galaxies \citep{vaughan}."274. The X-ray spectral properties of Z11598-0112. in particular Cie steep soft. X-ray spectrum: and Che flat haad X-ray spectrum. are consistent wilh a NLSI classification.," The X-ray spectral properties of Z11598-0112, in particular the steep soft X-ray spectrum and the flat hard X-ray spectrum, are consistent with a NLS1 classification."275 Twelve of the fourteen galaxies that we have observed with do not have enough counts for the usual spectral modeling procedure., Twelve of the fourteen galaxies that we have observed with do not have enough counts for the usual spectral modeling procedure.276 These sources have total counts in the keV band ranging from 3 to 92., These sources have total counts in the 0.5--8.0 keV band ranging from 3 to 92.277 In order to determine the properties of these sources. we used harcduess ratios to estimate modelparameters from XSPEC.," In order to determine the properties of these sources, we used hardness ratios to estimate modelparameters from XSPEC."278 The hardness ratio (1111), The hardness ratio (HR)279less computational burden than computing the entire integral in equation for each cell.,less computational burden than computing the entire integral in equation for each cell.280" For this comparison, we use the same initial disk as ?,, except that we rescale the physical parameters so that Mata;=1 Mo and Maia;=0.14 Mo."," For this comparison, we use the same initial disk as \cite{mejia2005}, except that we rescale the physical parameters so that $M_{\mathrm{star}} = 1$ $\mathrm{M_\odot}$ and $M_{\mathrm{disk}} = 0.14$ $\mathrm{M_\odot}$."281 The ? simulation results are also appropriately rescaled in all comparisons presented here., The \citeauthor{mejia2005} simulation results are also appropriately rescaled in all comparisons presented here.282 The disk extends from 2.3 to 40 AU with a surface density ©ος@~!/?., The disk extends from 2.3 to 40 AU with a surface density $\Sigma \propto \varpi^{-1/2}$.283" The computational gird has a central hole with a radius of 1.6 AU in which the star resides, and the 40 AU initial disk outer radius is at radial zone 242."," The computational gird has a central hole with a radius of 1.6 AU in which the star resides, and the 40 AU initial disk outer radius is at radial zone 242."284 The new run with the indirect potential is started from the same axisymmetric disk after applying the same 0.01% amplitude random cell-to-cell density perturbation used by ?.., The new run with the indirect potential is started from the same axisymmetric disk after applying the same $0.01\%$ amplitude random cell-to-cell density perturbation used by \citeauthor{mejia2005}.285" As in ?,, the constant £;55; is set equal to 2 Outer Rotation Periods or"," As in \citeauthor{mejia2005}, the constant $t_{\mathrm{cool}}$ is set equal to 2 Outer Rotation Periods or"286shifts (one bin = O.1 in redshitt} correspouding oz = 0.6. clesarily inelciting the prexnce of yeaks nn he distribution that are periodic z= 1.62.,"shifts (one bin = 0.1 in redshift) corresponding to z = $\pm0.6$, clearly indicating the presence of peaks in the distribution that are periodic in z = 0.62."287 That the highest value of r is obtained for a shif of -0.6 is not surprising since it corresponds o the situationi where the hielrest infriusic Dine deusitv (z  1.2) is aligued with the highest peak imtje quasar redshift disributicon (z 1.5)., That the highest value of r is obtained for a shift of -0.6 is not surprising since it corresponds to the situation where the highest intrinsic line density (z $\sim 1.2$ ) is aligned with the highest peak in the quasar redshift distribution (z $\sim 1.8$ ).288 This mrther confirms that the periodic peaks in the distribution below z = 2 are also coutributing to he correlation coefficieut., This further confirms that the periodic peaks in the distribution below z = 2 are also contributing to the correlation coefficient.289 The fact that x remains ugh over many shifts is due to he broad density correlation between the upper aid lower halves of he data., The fact that r remains high over many shifts is due to the broad density correlation between the upper and lower halves of the data.290 There was no significant correlation fonud vetween the SDSS redshift disribution and the distribution of lutrimsic redshifts predicted by equ 1., There was no significant correlation found between the SDSS redshift distribution and the distribution of intrinsic redshifts predicted by eqn 1.291" This is no surprising since. ia contrast o what was foud for equ 2. he iutiiusic lines Duicted Duπι d coincide wih oulv oue peak in tιο SDSS distribution (z-— 0.6). aud there is litle correlation between the source aud line densities in tle| upper and low""Y halves of the distributions."," This is not surprising since, in contrast to what was found for eqn 2, the intrinsic lines predicted by eqn 1 coincide with only one peak in the SDSS distribution (z = 0.6), and there is little correlation between the source and line densities in the upper and lower halves of the distributions."292 I therefore must |© coucluded that the model descrjbed by equ lis not coufiiuied by the SDSS redshMt data., It therefore must be concluded that the model described by eqn 1 is not confirmed by the SDSS redshift data.293 Furhonnore. While it has previously been demonstraed that the intrinsic redshifts predicted by equ 2 also show a siguificaut correlation wit[um oth he redshift distitiou of the 571 strong quasars Ua iuearly surveys (Bell2002¢:DellancComeau2003Pa ).. aud the redshift distribution of the QSOs located around NGC 6212 (BellandComet2010510). there was also no sienificant corrlation ound in these cases usine equ 1.," Furthermore, while it has previously been demonstrated that the intrinsic redshifts predicted by eqn 2 also show a significant correlation with both the redshift distribution of the 574 strong quasars found in early surveys \citep{bel02c,bel03b}, and the redshift distribution of the QSOs located around NGC 6212 \citep{bel03b}, there was also no significant correlation found in these cases using eqn 1."294 lu this exiuunmation of 11210 SDSS quasar redshifts it has ]ecu found that trere are at least 5 peaks in the recshift distibutioii that correspoud to the pure harmonics of z = (62 predicted. by equ 2.," In this examination of 44,200 SDSS quasar redshifts it has been found that there are at least 5 peaks in the redshift distribution that correspond to the pure harmonics of z = 0.62 predicted by eqn 2."295 Althoug1 valid arguineus lave been put forward to explain the redshif vallevs at zg = 2.7 and 3.5. theyo is no proof νεt that they have correctly explaired the causes ο these low source counts. siice eq 12 predicts low densities at both these redshifts. reducing the effect of any selection effect.," Although valid arguments have been put forward to explain the redshift valleys at z = 2.7 and 3.5, there is no proof yet that they have correctly explained the causes of these low source counts, since eqn 2 predicts low densities at both these redshifts, reducing the effect of any selection effect."296 Equ 2 aso predicts the low source count seen hear z = 03 where uo selection effect has vet been proposed., Eqn 2 also predicts the low source count seen near z = 4.03 where no selection effect has yet been proposed.297 It is concluded here that the iutrinsie redshift model given by equ 2 can better explain the structure present iu the distribution of these 11.200 SDSS quasar redshifts than can the selection effects sugeested by others.," It is concluded here that the intrinsic redshift model given by eqn 2 can better explain the structure present in the distribution of these 44,200 SDSS quasar redshifts than can the selection effects suggested by others."298 SDSS investigators are urged to keep this im ήτα if they hope to dusure that real density fluctuatious are not inadvertently removed., SDSS investigators are urged to keep this in mind if they hope to insure that real density fluctuations are not inadvertently removed.299 If we are to obtain a truly reliable statistical sample. adjustiueuts that affect the SDSS vedshift distribution mast not involve the assmuption that quasar redshifts are purely cosinological.," If we are to obtain a truly reliable statistical sample, adjustments that affect the SDSS redshift distribution must not involve the assumption that quasar redshifts are purely cosmological."300 Funding for the creation and distribution of the SDSS Archive has Όσοι. provided bw the Alfred P. Sloan Foundation. the Participating Tustiutious. the National Acronautics auc Space Aduunistration. the National Scieuce Foundation. the US Departinent of Encerex. the JapanOsc Moubukagalkusho. aud the Max Planck Society.," Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the US Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society."301 The SDSS website is lttp:/wwwesdss.ore/. The SDSS is managed by the Astrophysical Research Cousortimn (ARC) for the Participating Iustitutions., The SDSS website is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions.302 The Participating Iustitutious are Tιο University of Chicago. Fermilab. the Institute or Advance Study. The Japan Particiation Croup. The Johus Hopkins University. Los Alamos Natiolia Laboratory. the Max-Plauck-Institute for Astronomy (MIPTA). the Max-Plauck-Iusitute for Astroplisbes (AIPA). New Mexico State UAdversv. University of Pittsburgh. Princeton University. the Unite States Naval Observatory. aid the University of Washington.," The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, The Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington."303"The baryon budget shows significant evolution from z~3, and results in an apparent baryon deficit today (?7?)..","The baryon budget shows significant evolution from $z \sim 3$, and results in an apparent baryon deficit today \citep{FHP98,FP04}."304" At high redshift, most of the baryonic mass is in the Ly-o forest (??),, while at low redshifts over half of the baryons are as yet undetected."," At high redshift, most of the baryonic mass is in the $\alpha$ forest \citep{FHP98,FP04}, while at low redshifts over half of the baryons are as yet undetected."305" The consensus is that the majority of the *missing! baryons are actually in regions of low overdensity, range 10°—10"" K — commonly referred to as the Warm-Hot Intergalactic Medium (WHIM))."," The consensus is that the majority of the `missing' baryons are actually in regions of low overdensity, $\delta\rho/\rho \sim 10-100 $ \citep[e.g][]{CO99,DHKW99,DC01,KRCS05,CO06,DM06,DO07} with temperatures in the range $10^5-10^7$ K – commonly referred to as the Warm-Hot Intergalactic Medium )."306 The immediate question is: how is this produced?, The immediate question is: how is this produced?307 Some form of mass and energy injection is essential to create this hot reservoir of gas; this form of feedback must both regulate the gas in galaxies and the metal content of the Intergalactic Medium (IGM))., Some form of mass and energy injection is essential to create this hot reservoir of gas; this form of feedback must both regulate the gas in galaxies and the metal content of the Intergalactic Medium ).308" There has been much numerical work to incorporate various feedback mechanisms in an attempt to solve this puzzle (e.g., see"," There has been much numerical work to incorporate various feedback mechanisms in an attempt to solve this puzzle \citep[e.g.,][and references therein]{CO99,NS01,DC01,KRCS05,CO06,DO07}."309m to suggest that gravitational collapse during galaxy formation can produce and maintain the majority of the at 10?—107 K (?????)..," Cosmological simulations seem to suggest that gravitational collapse during galaxy formation can produce and maintain the majority of the at $10^5 - 10^7$ K \citep{CO99,DHKW99,CO06,CDM01,DO07}."310" Supernova feedback provides another avenue to generate the WHIM;; for star bursts of 100 Mg per year, as much as 20% of the hot gas in a Milky Way mass galaxy can be unbound (??).."," Supernova feedback provides another avenue to generate the ; for star bursts of 100 $M_\odot$ per year, as much as $20\%$ of the hot gas in a Milky Way mass galaxy can be unbound \citep{STWS06,KSW07}."311" However, SNe feedback may be a self regulating process, in that a starburst also heats the remaining gas and may damp the star formation rate, which in turn would reduce the fraction of unbound gas (e.g. ?).."," However, SNe feedback may be a self regulating process, in that a starburst also heats the remaining gas and may damp the star formation rate, which in turn would reduce the fraction of unbound gas \citep[e.g.][]{STWS08}. ."312" Combining these effects, it is commonly thought that galaxies with host halo mass >1011 Mo lose X10% of their gas through SN feedback, while low mass haloes may be entirely depleted of gas by this mechanism(???7).."," Combining these effects, it is commonly thought that galaxies with host halo mass $\gtrsim 10^{11} $ $_{\odot}$ lose $\lesssim 10\%$ of their gas through SN feedback, while low mass haloes may be entirely depleted of gas by this \citep{YK97,MF99,E00,STWS06}."313 A third possibility is that the radiation from an accreting supermassive black hole could power large-scale winds to blow mass out of the galaxy (see???7??7?7).," A third possibility is that the radiation from an accreting supermassive black hole could power large-scale winds to blow mass out of the galaxy \citep[see][]{SO04,MQT05,HH05,HHM05,HH06,CS06,SSMH07}."314". For a fixed amount of energy, all the non-gravitational feedback mechanisms are more effective in low mass galaxies due to their shallower potential."," For a fixed amount of energy, all the non-gravitational feedback mechanisms are more effective in low mass galaxies due to their shallower potential."315" However, observations suggest that low-mass galaxies are in general more gas-rich and are less likely to have suffered a gas (?7).."," However, observations suggest that low-mass galaxies are in general more gas-rich and are less likely to have suffered a gas blow-out \citep{K04,GBMW06}."316" In ?,hereafter,SH09,, we show that hot gas is driven into the by galaxy mergers."," In \citet[][hereafter, SH09]{SH09}, we show that hot gas is driven into the by galaxy mergers."317 The amount of hot halo gas lost depends strongly on the energy of the merger; it is possible for low mass galaxies to retain their gas in this scenario during low-speed or distant encounters., The amount of hot halo gas lost depends strongly on the energy of the merger; it is possible for low mass galaxies to retain their gas in this scenario during low-speed or distant encounters.318" However, SH09 only estimated the mass lost during a single merger."," However, SH09 only estimated the mass lost during a single merger."319" When allthe mergers inthe Universe are considered, this could heat and drive a significant portion of the total baryon"," When allthe mergers inthe Universe are considered, this could heat and drive a significant portion of the total baryon"320Iu such cases the analysis is carried through step 2. to determine the maxinmnmn response aud lence au estimate of the large separation.,"In such cases the analysis is carried through step 2, to determine the maximum response and hence an estimate of the large separation."321 Results ou the three individual cases are preseuted iu retsec:results.., Results on the three individual cases are presented in \\ref{sec:results}.322 Stellar evolution models aud aciabatic oscillation frequencies were computed using the Aarhus codes (Christensen-Dalseaard2008a.).. with the OPAL equation of state (Rogersetal.1996) and. opacity (lelesias&Rogers1990) and the NACRE— unclear reaction parameters (Aneuloetal.1999).," Stellar evolution models and adiabatic oscillation frequencies were computed using the Aarhus codes \citep{Christ2008a, Christ2008b}, with the OPAL equation of state \citep{Rogers1996} and opacity \citep{Iglesi1996} and the NACRE nuclear reaction parameters \citep{Angulo1999}."323. Tn some cases (see below) diffusion and settling of helium were included. using the simplified forxiiulation of MichaudProffitt (1993).," In some cases (see below) diffusion and settling of helium were included, using the simplified formulation of \citet{Michau1993}."324". Couvection was treated with the Bolun-Vitense(1958) uixine-leneth formulation. with a mixing leugth ayy,=2.00 iu uuits of the pressure scale height roughly correspouding to a solar calibration."," Convection was treated with the \citet{Bohm1958} mixing-length formulation, with a mixing length $\alpha_{\rm ML} = 2.00$ in units of the pressure scale height roughly corresponding to a solar calibration."325" Iu some models with convective cores. overshoot was included over a clistance of a, pressure scale heights."," In some models with convective cores, overshoot was included over a distance of $\alpha_{\rm ov}$ pressure scale heights."326 Evolution started from chemically homogencous zero-age models., Evolution started from chemically homogeneous zero-age models.327" The initial abundances by mass Ny and Zy of hydrogen and heavy eleiieuts were characterized by the assumed value of [Fe/TI]. using as reference a prescut solar surface composition with Z./.X,=0.0215 (Crevesse&Nocls1993) and assunuiug. from galactic chemical evolution. that Vy=(76793Zy."," The initial abundances by mass $X_0$ and $Z_0$ of hydrogen and heavy elements were characterized by the assumed value of [Fe/H], using as reference a present solar surface composition with $Z_{\rm s}/X_{\rm s} = 0.0245$ \citep{Greves1993} and assuming, from galactic chemical evolution, that $X_0 = 0.7679 - 3 Z_0$."328 From the observed Ary. effective temperature and composition an initial estimate of the stellar paraimcters was obtained using the erid-hased SEEN pipeline (Quirion et al.," From the observed $\Delta \nu_0$, effective temperature and composition an initial estimate of the stellar parameters was obtained using the grid-based SEEK pipeline (Quirion et al.,"329 in preparation)., in preparation).330 Smaller exids were then computed iu the vicinity of these initial paramicters. to obtain tighter constraints on stellaz properties.," Smaller grids were then computed in the vicinity of these initial parameters, to obtain tighter constraints on stellar properties."331" For ILAT-P-7 the analvsis of the observations yielded. frequeucies of individually identified modes: here the analysis was based ou where fol) ancl Uy(1iuod) iro the observed and model frequencies.os, 6; is the standard error im. the observed frequencies (asstuned to be coustaut) and Α΄ is the ummber of observed. frequencies."," For HAT-P-7 the analysis of the observations yielded frequencies of individually identified modes; here the analysis was based on where $\nu_{nl}^{\rm (obs)}$ and $\nu_{nl}^{\rm (mod)}$ are the observed and model frequencies, $\sigma_\nu$ is the standard error in the observed frequencies (assumed to be constant) and $N$ is the number of observed frequencies."332 Tn addition. we considered- qP=AZ|D where PEis the corespoudiug. jiorinalized square differenceyg betweenyp the observed aud nodel effective⋅⋅ temperature.," In addition, we considered $\chi^2 = \chi_\nu^2 + \chi_T^2$, where $\chi_T^2$ is the corresponding normalized square difference between the observed and model effective temperature."333" ⇁When yz» was available: we uiininized it along cach evolution track aud considered he resulting niüninmn values. aud the correspouding value of 47. as a function of the parameters characterizing he models (seeCullilandetal.201010,fordetails)."," When $\chi_\nu^2$ was available we minimized it along each evolution track and considered the resulting minimum values, and the corresponding value of $\chi^2$, as a function of the parameters characterizing the models \citep[see][for details]{Gillil2010b}."334.. When oulv the laree separation Amy could be determined from he observations. we identified the model along each track . : ↖↖↽↕∐↸⊳∐⋯⋜↧↑↸⊳↕∐∖≼↧∆↗∕⋃⋜⊔≼↧↸⊳∪∐↴∖↴↕≼∐∖↥⋅↸∖≼↧↑∐↸∖↥⋅↸∖↴∖↴∏↕⊓∐∶↴∙∖≣−⋜↧↴∖↴ ⋅⋝⋅≻ a fuuctiou of the iiodel parameters.," When only the large separation $\Delta \nu_0$ could be determined from the observations, we identified the model along each track which matched $\Delta \nu_0$ and considered the resulting $\chi_T^2$ as a function of the model parameters."335 The observed power spectrum for ILAT-P-7 is shown iu Fie., The observed power spectrum for HAT-P-7 is shown in Fig.336 laa. The presence of soku-like pauode peaks. with οΠΠ power around auallz. is evident.," \ref{fig:obs}a a. The presence of solar-like p-mode peaks, with a maximum power around mHz, is evident."337 At high frequency the noise level iu the amplitucle spectrum is l.l parts per million (ppii). with some increase at lower frequency. likely due to the effects of stellar eranulation.," At high frequency the noise level in the amplitude spectrum is 1.1 parts per million (ppm), with some increase at lower frequency, likely due to the effects of stellar granulation."338 Carrving out the correlation analvsis described iu refseciobs πο determined the large separation as Amy=59.22Iz.," Carrying out the correlation analysis described in \\ref{sec:obs}339 we determined the large separation as $\Delta \nu_0 = 59.22 \muHz$."340 Figure Hbb shows the resulting folded spectrmm., Figure \ref{fig:obs}b b shows the resulting folded spectrum.341 This clearly shows two closely spaced peaks. identified as corresponding to modes of degree |=0 and 2. and sinele peak separated from these two by approximately Amy/2. corresponding to /=1.," This clearly shows two closely spaced peaks, identified as corresponding to modes of degree $l = 0$ and 2, and single peak separated from these two by approximately $\Delta \nu_0/2$ , corresponding to $l = 1$."342" On this basis we finally determined the individual frequencies. identifving the modes from the asvurptotic relation: the final set ducludes 33 panode frequencies. determined with a standard eror a,τ.μ."," On this basis we finally determined the individual frequencies, identifying the modes from the asymptotic relation; the final set includes 33 p-mode frequencies, determined with a standard error $\sigma_\nu = 1.4 \muHz$."343 These frequencies.21. are illustrated in Fie.," These frequencies, are illustrated in Fig."344 Lee iun an écchelle diagram (see below)., \ref{fig:obs}c c in an écchelle diagram (see below).345 A erid of models was computed for masses between 1.11 aud 1.61M... |Fo/TI]. between 0.17. and 0.38. aud que0.0.1 and 0.2. extending well bevoud the end of ceutral hydrogen buruiug.," A grid of models was computed for masses between $1.41$ and $1.61 \,\Msun$, [Fe/H] between 0.17 and 0.38, and $\alpha_{\rm ov} = 0, 0.1$ and $0.2$, extending well beyond the end of central hydrogen burning."346 The modeling did not imclude diffusion aud settling., The modeling did not include diffusion and settling.347 At the mass of this star the outer convection zone is quite thin. aud as a result the settling nuescale is much shorter than the age of the star.," At the mass of this star the outer convection zone is quite thin, and as a result the settling timescale is much shorter than the age of the star."348 Tuchiding settling. without compcusating effects such as xwtial mixing iu the radiative region or nass loss. leads ο a vapid change in the surface composition which is inconsistent with the observed |Fe/TI]: for simplicity we herefore neelected these effects for The computed frequencies were corrected according to he procedure of Kjeldsenetal.(2008). for errors in the uodeline of the near-surface lavers. by adding ανν). where ¢=0.1158 plz. vy=1000Tz and b=L9.," Including settling, without compensating effects such as partial mixing in the radiative region or mass loss, leads to a rapid change in the surface composition which is inconsistent with the observed [Fe/H]; for simplicity we therefore neglected these effects for The computed frequencies were corrected according to the procedure of \citet{Kjelds2008} for errors in the modeling of the near-surface layers, by adding $a (\nu/\nu_0)^b$ where $a = 0.1158 \muHz$ , $\nu_0 = 1000 \muHz$ and $b = 4.9$."349 As discussed iui refsecimocdoel. for each evolution track. characterized by a set of model parameters; we minimized the departure \2 of the model frequencies from the observations. defining he best iiodel for this set.," As discussed in \\ref{sec:model}, for each evolution track, characterized by a set of model parameters, we minimized the departure $\chi_\nu^2$ of the model frequencies from the observations, defining the best model for this set."350 We first consider 4Z as a function of the effective eniperature of the models (Fig., We first consider $\chi_\nu^2$ as a function of the effective temperature of the models (Fig.351 2aa)., \ref{fig:chisq}a a).352 It is evident that lire is a clear diuinum in x2: this is consistent with the determination of Tig by Palctal.(2008) but not with he somewhat higher temperature obtained by. Anuuler-vonEifctal.(2009). (see also Table 1))., It is evident that there is a clear minimum in $\chi_\nu^2$; this is consistent with the determination of $T_{\rm eff}$ by \citet{Pal2008} but not with the somewhat higher temperature obtained by \citet{Ammler2009} (see also Table \ref{tbl-0}) ).353 Thus in the ollowiug we use the observed quautitics from Palctal. (2005)., Thus in the following we use the observed quantities from \citet{Pal2008}.354 Since the frequencics to leading order are determined w the mean stellar density £9.). Fig.," Since the frequencies to leading order are determined by the mean stellar density $\rhomean$, Fig."355 2bb.c show 4Z aud (2 as functions of ο," \ref{fig:chisq}b b,c show $\chi_\nu^2$ and $\chi^2$ as functions of $\rhomean$."356 Tt is evident that the best-ftting nodels occupy a narrow rauge of (9.5. with a well-defined uiniunun.," It is evident that the best-fitting models occupy a narrow range of $\rhomean$, with a well-defined minimum."357 Fitting a parabola to 4? in panel (c) we obtain he estimate £p.)=0.2712£0.0032¢cm!.," Fitting a parabola to $\chi^2$ in panel (c) we obtain the estimate $\rhomean = 0.2712 \pm 0.0032 \,{\rm g \, cm^{-1}}$."358 In Fig., In Fig.359 2dd (P is shown against model age., \ref{fig:chisq}d d $\chi^2$ is shown against model age.360 Tere the variation with nodel paraicters is substantially stronger. resulting iu a ereater spread in the inferred) age: in particular. it is evident. not surprisiuglw. that the results depend ou he extent of convective overshoot.," Here the variation with model parameters is substantially stronger, resulting in a greater spread in the inferred age; in particular, it is evident, not surprisingly, that the results depend on the extent of convective overshoot."361 From the figure weestimatethat the age of ITAT-P-7 is between L1 aud GC., From the figure weestimatethat the age of HAT-P-7 is between 1.4 and Gyr.362 Examples of evolution tracks are shown in Fig. 3: , Examples of evolution tracks are shown in Fig. \ref{fig:HR}; ;363piraueters for these models are provided iu Table 2.., parameters for these models are provided in Table \ref{tbl-2}. .364Our derivation also allows us to calculate other quantities related to the distribution of strains.,Our derivation also allows us to calculate other quantities related to the distribution of strains.365" First. we can ask what is the ummber of binaries in a frequency interval. Using equation (27)). and again under the assuniptious we have made so far that the o, is independent of redshift. the + integration separates out. and we sce first that ΑΕ)df/foxfοσαf. a steeply-falliug function of frequency."," First, we can ask what is the number of binaries in a frequency interval, Using equation \ref{eq:strain1z}) ), and again under the assumptions we have made so far that the $\phi_\BH$ is independent of redshift, the $z$ integration separates out, and we see first that $N(f)\; df/f \propto366f^{-8/3} df/f$, a steeply-falling function of frequency."367 In. full. CMCDS.-m ds ⋅∣↽↽(LIsLOPS.) RNfor our MDBII mass fuuctiou. which⋅⋅⋅⋅ is different from. the value above owing to a different weighting fuuction.," In full, $\langle{\cal M}^{-5/3}\rangle_{\BH}$ is $\left(4.1 \times 10^6368 M_\odot\right)^{-5/3}$ for our MBH mass function which is different from the value above owing to a different weighting function."369 In Fieure 5 we show the dimensionless iuteeral occurring iu this expression for a variety of cosnoloeios aud merger bistories., In Figure \ref{fig:Nf} we show the dimensionless integral occurring in this expression for a variety of cosmologies and merger histories.370 The stochastic GW background at ullz (vr1) frequencies is therefore the result of nearly a οι simultaneous binaries across the Universe., The stochastic GW background at nHz $^{-1}$ ) frequencies is therefore the result of nearly a million simultaneous binaries across the Universe.371"results are Again, our result for 0, differs from that given in[KC84] The results for an unmagnetized relativistic shock are recovered by simply setting c=0 in the above relations.","results are Again, our result for $\theta_d$ differs from that given in The results for an unmagnetized relativistic shock are recovered by simply setting $\sigma=0$ in the above relations."372" Here we consider the non-relativistic case and replace u,,ug by Bu, Ba."," Here we consider the non-relativistic case and replace $u_u, ~u_d$ by $\beta_u, ~\beta_d$ ."373" Also, we set h(0,)=5/2."," Also, we set $h(\theta_d)=5/2$."374 Then we find The solution for an unmagnetized shock is obtained by setting o=0., Then we find The solution for an unmagnetized shock is obtained by setting $\sigma=0$.375" As an aside, we note that o is related to the Alfven wave speed v4 by Finally, we consider the case when the shock is non-relativistic and By=Na(1+A) with Vo«A1."," As an aside, we note that $\sigma$ is related to the Alfven wave speed $v_A$ by Finally, we consider the case when the shock is non-relativistic and $\beta_u = \sqrt{\sigma}(1+\Delta)$ with $\sqrt\sigma\ll\Delta \ll 1$."376" In this limit, we find We generalize the discussion of the parameter £ given in[SP95], following the analysis of"," In this limit, we find We generalize the discussion of the parameter $\xi$ given in, following the analysis of."377" We consider a spherically expanding shell of cold [Gianniosetmagnetizedal] (2008)..jet material of radius R, shell thickness A, and Lorentz factor y;, all measured in the lab frame."," We consider a spherically expanding shell of cold magnetized jet material of radius $R$, shell thickness $\Delta$, and Lorentz factor $\gamma_j$, all measured in the lab frame."378" The “spreading radius"" of the shell is given by The total (isotropic equivalent) energy of the shell is where ny is the rest frame particle number density of the jet material, o is the magnetization of the material, and M,;j is the total rest mass of the shell."," The “spreading radius” of the shell is given by The total (isotropic equivalent) energy of the shell is where $n_4$ is the rest frame particle number density of the jet material, $\sigma$ is the magnetization of the material, and $M_{\rm379ej}$ is the total rest mass of the shell."380" From £, we obtain the “Sedov length"" £ and the “deceleration radius"" Race, where πι is the number density of the external ambient medium."," From $E$, we obtain the “Sedov length” $\ell$ and the “deceleration radius” $R_{\rm dec}$, where $n_1$ is the number density of the external ambient medium."381" Substituting for Mj (with R=Rgec) in the equation for Réec, We find thatfrom which we obtain Note thatn4 is the number density of the jet ejecta at the moment when the shell radius R is equal to Raec."," Substituting for $M_{\rm ej}$ (with $R=R_{\rm dec}$ ) in the equation for $R_{\rm dec}$, we find thatfrom which we obtain Note that$n_4$ is the number density of the jet ejecta at the moment when the shell radius $R$ is equal to $R_{\rm dec}$ ."382on the 2ALASS svstem.,on the 2MASS system.383 Iteassuringlv. these results cülfer from those derived. using our large parallax! subset. (see above) by only 6-Smmag.," Reassuringly, these results differ from those derived using our 'large parallax' subset (see above) by only 6-8mmag."384 Likewise reassuring is the fact that a comparison of these results with those including only the nearer stars shows no sign whatever of extinction elfects., Likewise reassuring is the fact that a comparison of these results with those including only the nearer stars shows no sign whatever of extinction effects.385 Given past interest in the effect of metal abundance on red clump absolute magnitudes. we examined our data to see if any trend was apparent.," Given past interest in the effect of metal abundance on red clump absolute magnitudes, we examined our data to see if any trend was apparent."386 From our sample. 101 stars had metal abunclances either from AleWilliam (1990) or Liu et al. (," From our sample, 101 stars had metal abundances either from McWilliam (1990) or Liu et al. ("3872007).,2007).388 A comparson of 24 stars in common showed that the abundances [rom these two sources had. cdillerent zero points. and that. the AleWilliam metallicities could. be placed. on the scale. of Liu et al.," A comparson of 24 stars in common showed that the abundances from these two sources had different zero points, and that the McWilliam metallicities could be placed on the scale of Liu et al."389 simply by aciing O.12+0.02., simply by adding $\pm$ 0.02.390 For the stars with abundances from both sources. t1ο two values have been averaged.," For the stars with abundances from both sources, the two values have been averaged."391 Absolute magnitudes in. ly.4piss and Llosaiss have been plotted against metallicity (on the Liu et al., Absolute magnitudes in $_{2MASS}$ and $_{2MASS}$ have been plotted against metallicity (on the Liu et al.392 scale) in Figs., scale) in Figs.393 3 and 4., 3 and 4.394 As is evident fron1 these figures. there is no strong. significant trend in either absolute magnitude with metallicity (at least for stars with MILI] greater than -0.6). in agreement with the result found by Alves (2000).," As is evident from these figures, there is no strong, significant trend in either absolute magnitude with metallicity (at least for stars with [M/H] greater than -0.6), in agreement with the result found by Alves (2000)."395 To derive an LMC. distance. modulus. we need a mean lxouass value for LAIC red clump stars.," To derive an LMC distance modulus, we need a mean $_{2MASS}$ value for LMC red clump stars."396 “Phe two most comprehensive studies of LMC red clump stars in the field eive dereddened mean Isa;àss magnitudes of 550.000 (Alves ct al., The two most comprehensive studies of LMC red clump stars in the field give dereddened mean $_{2MASS}$ magnitudes of $\pm$ 0.009 (Alves et al.397 2002) ancl 16.89740.009. (Szewezvk ct al., 2002) and $\pm$ 0.009 (Szewczyk et al.398 2008). while an extensive recent survey of LMC red clump stars in clusters (Grocholski et al.," 2008), while an extensive recent survey of LMC red clump stars in clusters (Grocholski et al."399 2007) vields a mean of 1G.S91+0.082., 2007) yields a mean of $\pm$ 0.032.400 Phese are in extremely good agreement. and have been averaged to give We have followed Grocholski ct al. (," These are in extremely good agreement, and have been averaged to give We have followed Grocholski et al. ("4012007) in using he reddening law of Cardelli et al. (,2007) in using the reddening law of Cardelli et al. (4021989).,1989).403 The mean Ix magnitudes from Alves e al. (, The mean K magnitudes from Alves et al. (4042002) and Grocholski et al. (,2002) and Grocholski et al. (4052007) have been corrected to the LMC center as described in those papers. while the mean Ix magnitude from Szewezvk et al. (,"2007) have been corrected to the LMC center as described in those papers, while the mean K magnitude from Szewczyk et al. ("4062008) has been left uncorrected for reasons cited by he authors of that paper.,2008) has been left uncorrected for reasons cited by the authors of that paper.407 In all cases the Ix magnitudes from hese three papers have been transformed. to the ολλ system using the transformations on the 2ALASS website (Carpenter 2003)., In all cases the K magnitudes from these three papers have been transformed to the 2MASS system using the transformations on the 2MASS website (Carpenter 2003).408 The mean Ix. magnitudes for red. clump, The mean K magnitudes for red clump409in different phases of the modulation.,in different phases of the modulation.410 The animated light curves of the M5 Blazhko variables and the data sets that were used are available at the url::www., The animated light curves of the M5 Blazhko variables and the data sets that were used are available at the url:.411konkoly.hu/24/publications/M5. We identified 20 Blazhko candidates among the 50 RRab variables that were studied., We identified 20 Blazhko candidates among the 50 RRab variables that were studied.412 The properties of the individual stars are discussed below., The properties of the individual stars are discussed below.413 There is a slight difference between the shape of minima observed by R96 and K00., There is a slight difference between the shape of minima observed by R96 and K00.414 The maximum brightness in the R96 data varies between 14.41 and 14.45-mag., The maximum brightness in the R96 data varies between 14.41 and 14.45-mag.415" The star is far from the centre, so it can be measured accurately on the photographic plates."," The star is far from the centre, so it can be measured accurately on the photographic plates."416 The residual spectrum of the combined photographic data between JD 2441447 and JD 2449104 (1972-1993) shows a large-amplitude (0.05-mag) signal at fo+0.001916 cd!., The residual spectrum of the combined photographic data between JD 2441447 and JD 2449104 (1972–1993) shows a large-amplitude (0.05-mag) signal at $f_0+0.001916$ $^{-1}$.417 We found 522 d as the probable period of the modulation for this time interval., We found 522 d as the probable period of the modulation for this time interval.418 The modulation is dominated by &20 min phase oscillations., The modulation is dominated by $\approx20$ min phase oscillations.419 'This result is supported by the analysis of the variation of the seasonal O—C' values of the entire photometric data (1935-1997) as shown in Fig. 1.., This result is supported by the analysis of the variation of the seasonal $O-C$ values of the entire photometric data (1935–1997) as shown in Fig. \ref{v01oc}.420 Phase oscillations with a period of 512 d fit the observed 0.02-d variation of the seasonal O—C data reasonably well., Phase oscillations with a period of 512 d fit the observed 0.02-d variation of the seasonal $O-C$ data reasonably well.421 T'he CCD data do not contradict this modulation period., The CCD data do not contradict this modulation period.422 The pulsation period of V1 remained constant during the time base of more than 100 years of the observations., The pulsation period of V1 remained constant during the time base of more than 100 years of the observations.423 The star is well separated., The star is well separated.424" The pulsation period did not show any systematic changes during the hundred years of the observations, but irregular period changes of the order of 107? d were detected."," The pulsation period did not show any systematic changes during the hundred years of the observations, but irregular period changes of the order of $10^{-5}$ d were detected."425 CCD data from K00 show different descending branches., CCD data from K00 show different descending branches.426 Already Oo41 noted differences between two maxima., Already Oo41 noted differences between two maxima.427 A modulation period of 132.38 d was determined by Goranskij(1976).., A modulation period of 132.38 d was determined by \cite{g76}.428 We have found clear evidence of the modulation in three segments of the combined photographic data., We have found clear evidence of the modulation in three segments of the combined photographic data.429" Between 1971 and 1986, V2 exhibited a very strong modulation with a 136.5-d period."," Between 1971 and 1986, V2 exhibited a very strong modulation with a 136.5-d period."430" The pulsation amplitude varied between 0.6 and 1.3-mag; the detected changes in the maximum light were 0.6-mag and 75-min in brightness and phase, respectively."," The pulsation amplitude varied between 0.6 and 1.3-mag; the detected changes in the maximum light were 0.6-mag and 75-min in brightness and phase, respectively."431"axtrometric accuracy of the catalogue is about 0.2"" at J2000.",astrometric accuracy of the catalogue is about $^{\prime\prime}$ at J2000.432 About fifteen USNO-D1.0 stars were found witLin the OSIS FOV around Toll X-1., About fifteen USNO-B1.0 stars were found within the OSIS FOV around HoII X-1.433" The standard deviations of the star positious are 0.20% for the Πα, 0.13” for the D. aud O15"" for the A images."," The standard deviations of the star positions are $^{\prime\prime}$ for the $H\alpha$ , $^{\prime\prime}$ for the $B$ , and 0.15"" for the $R$ images."434 In Fig., In Fig.435 2 we present the CFITT images with the ACIS-S source position marked (see detailed description in Sect., \ref{CFHT} we present the CFHT images with the ACIS-S source position marked (see detailed description in Sect.436 5)., 5).437 An object ou the D image perfectly coincides with the position of the Xταν source., An object on the $B$ image perfectly coincides with the position of the X–ray source.438 Its maguitude is 2=20.540.1 mae., Its magnitude is $B=20.5\pm0.1$ mag.439 Since all the stars on the D nuage in Fig., Since all the stars on the $B$ image in Fig.440 2 look elougated. while the seeing was fairly good (zx0.77). we fitted 2d-Gaussiaus to 23 stars surrounding he rav position.," \ref{CFHT} look elongated, while the seeing was fairly good $\approx 0.7$ ""), we fitted 2d-Gaussians to 23 stars surrounding the X-ray position."441" The elongation of these stars was big=1.36 40.05, where the major axis size is b=0.894E 0.017 aud the minor axis is ¢=0.65 0.017."," The elongation of these stars was $b/a = 1.36 \pm 0.05$ , where the major axis size is $b = 0.89 \pm 0.04$ "" and the minor axis is $a = 0.65 \pm 0.01$ ""."442 The position anele of the star mages was P.A.=19.2$:30., The position angle of the star images was $=-19.2 \pm 3.0^{\circ}$.443 The object coincident with the N-rav source has the following parzuneters: bfazz1.20. 5= 1.227. 022 0.967 and PASLO. which indicates that in comparison to the surrounding stars the counterpart is au extended object. e.g. a nebula or a compact stellar cluster.," The object coincident with the X-ray source has the following parameters: $b/a \approx 1.29$, $b \approx 1.22$ "", $a \approx 0.96$ "" and $ \approx -10^{\circ}$, which indicates that in comparison to the surrounding stars the counterpart is an extended object, e.g. a nebula or a compact stellar cluster."444" Its intrinsic size is about 1.69"" (11 pe) in the West-East direction and about 0.85"" (13 pc) in the North-South direction.", Its intrinsic size is about $^{\prime\prime}$ (11 pc) in the West-East direction and about $^{\prime\prime}$ (13 pc) in the North-South direction.445 The orientation is (07., The orientation is $\sim 0^{\circ}$.446 Assuming the B-baud flux calibrations frou: Allen (1973)). and a line-ofsieht extinction to ΠΟΠ of Ap=0.1I (Schlegel et al. 1998)).," Assuming the $B$ -band flux calibrations from Allen \cite{All73}) ), and a line-of-sight extinction to HoII of $A_B = 0.14^m$ (Schlegel et al. \cite{Schl98}) ),"447 the optical luminosity of the counterpart is estimated to ο Lg~6+107 org/s. The absolute magnitude of this object is Mg=7.2.," the optical luminosity of the counterpart is estimated to be $L_B \sim 6 \cdot44810^{37}$ erg/s. The absolute magnitude of this object is $M_B = -7.2$."449" Using the X-rav luminosity in the 0.3-8.0 keV energy baud of L,~1079 cre/s. corrected for absorption (see Sect."," Using the X-ray luminosity in the 0.3-8.0 keV energy band of $L_x \sim 10^{40}$ erg/s, corrected for absorption (see Sect."450 5.2). we fiud L./Lp7170.," 5.2), we find $L_x/L_B \ge 170$."451 ReceutLST ACS observations of Πο N-1. (published after the stbinission of this paper). resolved this object iuto several voung stars (INaaret ct al. 2001).," Recent ACS observations of HoII X-1, (published after the submission of this paper), resolved this object into several young stars (Kaaret et al. \cite{Kaa04}) ),"452 which agrees with our interpretation., which agrees with our interpretation.453 Due to the superb aueular resolution of the ACS images Riuuet found a bright. poiut-like optical counterpart consisteut either with a star with spectral type between OVand B3 Th. or reprocessed enission from au X-ray illuminated accretion disk.," Due to the superb angular resolution of the ACS images Kaaret found a bright, point-like optical counterpart consistent either with a star with spectral type between O4Vand B3 Ib, or reprocessed emission from an X-ray illuminated accretion disk."454 Because the star fouud by Naaret et al. (2001)), Because the star found by Kaaret et al. \cite{Kaa04}) )455" is about one maguitude fainter compared o our extended blue counterpart this results ii a L,./Lp of 300—100.", is about one magnitude fainter compared to our extended blue counterpart this results in a $L_x/L_B$ of $\sim300-400$.456 The MPFS Ue II Al6Gs6 line fix map (sce Fig., The MPFS He II $\lambda4686$ line flux map (see Fig.457 6 in Sect., \ref{mpfs_flux1} in Sect.458 L1) clearly shows a He IT ciission line region at the ACIS-S position. and coincident with the optical counterpart detected on the CHET B image.," 4.1) clearly shows a He II emission line region at the ACIS-S position, and coincident with the optical counterpart detected on the CHFT $B$ image."459 This confirms the cassification of ULX Πο X-1 as an X-rav domized nebula (Pakull Mirioni 20023)., This confirms the classification of ULX HoII X-1 as an X-ray ionized nebula (Pakull Mirioni \cite{Pak01}) ).460 Iu order to determine if the Te ID emission is extended we have derived the surface brghtuess profile of the Πο II region. using an algorithm based ou JedrzejewskEki (1987)).," In order to determine if the He II emission is extended we have derived the surface brightness profile of the He II region, using an algorithm based on Jedrzejewski \cite{Jed87}) )."461 This method increases the signal-to-noise in the outer part of the surface brightuess profile. since it comprises au average of the brigltuess alone the eccentric anomaly.," This method increases the signal-to-noise in the outer part of the surface brightness profile, since it comprises an average of the brightness along the eccentric anomaly."462 A simular technique is extensively used for the detection of ιοί ealaxies in QSOs (eg..," A similar technique is extensively used for the detection of host galaxies in QSOs (eg.,"463 Sáunchez Gouzalez-Serrauo 2003))., Sánnchez Gonzalez-Serrano \cite{San03}) ).464 The MPFS PSF was lilt using the /contmuun chussion at a wavoleugth range near the We II euission., The MPFS PSF was built using the continuum emission at a wavelength range near the He II emission.465 Fie., Fig.466 3 shows the surface brightuess profile of the He II region together with the surface brightuess profile of the PSF. scaled to the peak of the Te IT eimissiou.," \ref{profile}467 shows the surface brightness profile of the He II region together with the surface brightness profile of the PSF, scaled to the peak of the He II emission."468" The Πο II cluission is clearly extended bevoud τη, which is also confirmed by the ΙΤ ACS Πο ID narrow band image of TIkaaret et al. (2001))."," The He II emission is clearly extended beyond $\sim$ $^{\prime\prime}$ , which is also confirmed by the ACS He II narrow band image of Kaaret et al. \cite{Kaa04}) )."469" We find an extended aud elougated We II A£686 region with nearly the same positional angle of P.A.~19"" as found for the blue counterpart on the B image.", We find an extended and elongated He II $\lambda4686$ region with nearly the same positional angle of $\sim-19^{\circ}$ as found for the blue counterpart on the $B$ image.470 The size of the Πο II ALGS6 reeion after PSF correction is about l.l « 3.0% (see Fie. 3)).," The size of the He II $\lambda4686$ region after PSF correction is about 1.4 $\times$ $^{\prime\prime}$ (see Fig. \ref{profile}) ),"471 which corresponds to about 21 * LF pe., which corresponds to about 21 $\times$ 47 pc.472 The avecr size of the He ITI region compared with the blue counterpart «11 pc) could be considered as an argument for a stellar complex. where the D-baud counterpart represents the coutimmui emission aud the Ie IIl region represents the ligh excitation nebula.," The larger size of the He II region compared with the blue counterpart $\times$ 14 pc) could be considered as an argument for a stellar complex, where the $B$ -band counterpart represents the continuum emission and the He II region represents the high excitation nebula."473 The most miportaunt observational results from this section are that the Πο TT emission is: extended. aud in the sale positional auele as the extended blue counterpart. aud centered on the N-rayv source.," The most important observational results from this section are that the He II emission is: extended, and in the same positional angle as the extended blue counterpart, and centered on the X-ray source."474 This coufiris the classification of ΠΟΠ N-1 as an N-ray ionized rebula as sugeested by. (Palzul Mirioni 2002))., This confirms the classification of HoII X-1 as an X-ray ionized nebula as suggested by (Pakull Mirioni \cite{Pak01}) ).475 To determine the plysical parameters (e.g. the radial velocities and the velocity dispersious) of the IT IT regiou associated with the ULX. we have measured the emission line propertiesofall lines in the LSS. PMÁS aud \IPFS spectra.," To determine the physical parameters (e.g. the radial velocities and the velocity dispersions) of the H II region associated with the ULX, we have measured the emission line propertiesofall lines in the LSS, PMAS and MPFS spectra."476 Whereas the PALAS spectra cover ouly the Ie IT A L686. ITJASGT and the [O TH) AA1959.5007 cuiission lines. theMIPFS spectralrange includes the T+A1363 aud TaA6563 cunission line regious as well.," Whereas the PMAS spectra cover only the He II $\lambda4686$ , $\beta~\lambda 4861$ and the [O III] $\lambda\lambda 4959, 5007$ emission lines, theMPFS spectralrange includes the $\gamma~\lambda4363$ and $\alpha~\lambda6563$ emission line regions as well."477 Each cinission liue iu the 210 MPFES aud 506 PALAS integral field spectra was fitted with a Gaussian. profile, Each emission line in the 240 MPFS and 506 PMAS integral field spectra was fitted with a Gaussian profile478uegative polarity solutions have different properties: For the same values of all the other parameters. both the base of the wind aud the sonic point are located higher above the midplane when the magnetic ficld is aligned with the aneular velocity vector of the disc (B.> 0) than when it is pointing in the other direction (B. 0).,"negative polarity solutions have different properties: For the same values of all the other parameters, both the base of the wind and the sonic point are located higher above the midplane when the magnetic field is aligned with the angular velocity vector of the disc $B_z > 0$ ) than when it is pointing in the other direction $B_z < 0$ )."479 The mass outflow rate is correspondinglv lower in the former case than iu the latter., The mass outflow rate is correspondingly lower in the former case than in the latter.480" Furthermore. no viable solutiois are predicted to exist in these sub-regimes when Ty<2|iy)+ and B,<0."," Furthermore, no viable solutions are predicted to exist in these sub-regimes when $\upo < 2 \betiabs^{-1}$ and $B_z < 0$."481 For the other two parameter sub-regiues (Cases 7 aud fe du Table 1). the situation is even more drastic: Solutions are predicted to exist for the positive polarity only (c.g. πο viable solutions are expected for the negative polarity).," For the other two parameter sub-regimes (Cases $ii$ and $iv$ in Table \ref{table:constraints}) ), the situation is even more drastic: Solutions are predicted to exist for the positive polarity only (e.g. no viable solutions are expected for the negative polarity)."482 Our numerical solutious coufinii the above expectations, Our numerical solutions confirm the above expectations.483 This is illustrated iu Fig., This is illustrated in Fig.484 6— (for sub-reguue 7)., \ref{fig:5_26} (for sub-regime $i$ ).485" This figure shows the normalised deusity at the sonic point (p,/pog) as a function of the parameter Ty. for solutions with both (positive aud negativo) magnetic field polarities."," This figure shows the normalised density at the sonic point $\rho_{\rm s}/\rho_{\rm 0}$ ) as a function of the parameter $\upo$, for solutions with both (positive and negative) magnetic field polarities."486 Tn all cases. the value of |4|.1=1.15.," In all cases, the value of $\betiabs^{-1} = 1.45$."487 The displayed results confirm the prediction of the livdrostatic analysis (IRSW10) that πο viable solutions exist for ni—guO/e«0 when Ty5! (ic. to the left of the vertical dashed Hue)., The displayed results confirm the prediction of the hydrostatic analysis (KSW10) that no viable solutions exist for $\tilde{\eta}_{\rm H} \equiv \eta_{\rm H} \Omega/ c_{\rm s}^2 < 0$ when $\upo < 2 \betiabs^{-1}$ (i.e. to the left of the vertical dashed line).488" It is also evident that iu the region of parameter space where solutions with both polarities exist. these have differeut properties,"," It is also evident that in the region of parameter space where solutions with both polarities exist, these have different properties."489" Note that the mass outflow rate (1ieasured by psf py) is sinaller in the D,>O case, In accordance with our predictions."," Note that the mass outflow rate (measured by $\rho_{\rm s}/\rho_{\rm 0}$ ) is smaller in the $B_z > 0$ case, in accordance with our predictions."490 The AMIRI (Balbus&IawleyL991. see also the review by Balbus&Tawley 1998)). transfers angular momentum racially outwards via magnetic field. lines that connect fluid clemenuts located at different radii.," The MRI \citealt{BH91}, see also the review by \citealt{BH98}) ), transfers angular momentum radially outwards via magnetic field lines that connect fluid elements located at different radii."491 The physical principle that wuderpins this mechanisin is illustrated in Fie. 7.., The physical principle that underpins this mechanism is illustrated in Fig. \ref{fig:MRIcartoon}.492 huagine two fluid clemeuts that are initially located at the same orbital radius aud are joined bv a vertical magnetic field line threading the disc (eft panel)., Imagine two fluid elements that are initially located at the same orbital radius and are joined by a vertical magnetic field line threading the disc (left panel).493 In this configuration the maguetic tension in the line is negligible., In this configuration the magnetic tension in the line is negligible.494 Now imagine that the eleiieuts are perturbed from their initial orbits. so that oue of thei (the one labelled iy in the figure) is displaced to au inrer orbit. whereas the other (05) moves to an outer orbit (middle panel}.," Now imagine that the elements are perturbed from their initial orbits, so that one of them (the one labelled $m_{\rm 1}$ in the figure) is displaced to an inner orbit, whereas the other $m_{\rm 2}$ ) moves to an outer orbit (middle panel)."495 Asstming that the field-matter coupling is sufficient for the maguetic field line to be dragecd by the moving fluid elemieuts. the line will be stretched and. magnetic tension will develop.," Assuming that the field-matter coupling is sufficient for the magnetic field line to be dragged by the moving fluid elements, the line will be stretched and magnetic tension will develop."496 Note that the eleiieut in the iuner orbit (204) dMbvertakes the oue in the outer orbit (53). as a result of the depeudence of the Keplerian orbital inotion with radius (ey=VGM/r).," Note that the element in the inner orbit $m_{\rm 1}$ ) overtakes the one in the outer orbit $m_{\rm 2}$ ), as a result of the dependence of the Keplerian orbital motion with radius $v_{\rm K} = \sqrt{GM/r}$ )."497 The tension in the line. therefore. acts to reduce the aneular momentum of the muer clemeut. and to increase that of the outer clement (rieht pauel)h," The tension in the line, therefore, acts to reduce the angular momentum of the inner element, and to increase that of the outer element (right panel)."498 As a result of this process. aneulu moment is transferred from my to gn». causing the former clement to move further inwards while the latter retreats to au even larger radius.," As a result of this process, angular momentum is transferred from $m_{\rm 1}$ to $m_{\rm 2}$, causing the former element to move further inwards while the latter retreats to an even larger radius."499" As the elements continue to separate. the teusion in the magnetic field line joimime them imereases. and the process ""r1us away: perturbations. once initiated. will auuplify."," As the elements continue to separate, the tension in the magnetic field line joining them increases, and the process “runs away""; perturbations, once initiated, will amplify."500 The noulnear stages of the instability have been studied exeusivelv. (6.9.seethereviews|wBalbus&Tawley20025. 1998).," The nonlinear stages of the instability have been studied extensively \citep[e.g. see the reviews by][]{balbus03a, BH98}."501. The resut is MIID-diveu urbuleuce where angular momenttni ds trallspored radially ouwards., The result is MHD-driven turbulence where angular momentum is transported radially outwards.502 This outward trasport is linked to he fact that the turbulence is iixtropic: the racial aud aziauthal components of the pcvturbatious of he naenetic axd velocity fields are siguificatlv ccorrelated., This outward transport is linked to the fact that the turbulence is anisotropic: the radial and azimuthal components of the perturbations of the magnetic and velocity fields are significantly correlated.503 Tn other words. the transport of ajeular monientimm ollows directly from the statistical correlation iu he," In other words, the transport of angular momentum follows directly from the statistical correlation in the"504Finally. Li production would require temperatures not exceeding ~5x10* Ix in the processing zone.,"Finally, Li production would require temperatures not exceeding $\sim 5\times50510^7$ K in the processing zone."506 At least the last temperature requirement appear rather difficult to attain (Lattanzio. private communication).," At least the last temperature requirement appear rather difficult to attain (Lattanzio, private communication)."507 Perhaps. Li can be produced after the Ie-flash by. an additional mixing event between the convective envelope and the I-shell.," Perhaps, Li can be produced after the He-flash by an additional mixing event between the convective envelope and the H-shell."508 Obviously. the lle-flash scenario merits further hvdrodyvnamie (3-D) studies.," Obviously, the He-flash scenario merits further hydrodynamic (3-D) studies."509 Note. the Ie-flash occurs in low-mass and low-Iuminositv objects such the stars studied here.," Note, the He-flash occurs in low-mass and low-luminosity objects such the stars studied here."510 Finally. we consider the mass-transfer scenario in a binary svstem.," Finally, we consider the mass-transfer scenario in a binary system."511 The existence of 5 stars will no Te. as predicted bv (the mass-transfer paracien. is now well established (Jorissen.FraverandJohnson.1993).," The existence of S stars with no Tc, as predicted by the mass-transfer paradigm, is now well established \citep{jor93}."512. Whether this scenario can also be applied (to C-stars is not. vel firmly demonstrated. although. Darnbaunm(1993). found 16 Te-poor stars in a sample of 78 C-N stars with Ba excess.," Whether this scenario can also be applied to C-stars is not yet firmly demonstrated, although \citet{bar93} found 16 Tc-poor stars in a sample of 78 C-N stars with Ba excess."513 However. it is difficult to explain the absence of s-process element enhancement and (he C/O ratios in our stars within this scenario.," However, it is difficult to explain the absence of s-process element enhancement and the C/O ratios in our stars within this scenario."514 In principle. the accreted material must be extremely. carbou-rich: (he donor star should be a normal C-star with probably enhanced s-nuclei in the envelope.," In principle, the accreted material must be extremely carbon-rich; the donor star should be a normal C-star with probably enhanced s-nuclei in the envelope."515 Dominy(1985). estimated a C/O5 in the material accreted by alypical 1 AL. red giant when applving this scenario to explain the (051 ratios observed in R-stars., \citet{dom85} estimated a $\gtrsim 5$ in the material accreted by a $\sim 1$ $_\odot$ red giant when applying this scenario to explain the $>1$ ratios observed in R-stars.516 The same figure can be applied in the case of J-stars., The same figure can be applied in the case of J-stars.517 This extreme C/O ratio is not observed in any Costar., This extreme C/O ratio is not observed in any C-star.518 Furthermore. even assuming that (he material (rauslerrecl were Li-rich (some N-(vpe carbon stars are Li-rich). it is unlikelv (hat Li could survive during the mass-(ransler ancl posterior mixing.," Furthermore, even assuming that the material transferred were Li-rich (some N-type carbon stars are Li-rich), it is unlikely that Li could survive during the mass-transfer and posterior mixing."519 In fact. extrinsic (binary. no Te) 9 stars do not usually show the Li enhancements found here (Barbuyetal.1992).," In fact, extrinsic (binary, no Tc) S stars do not usually show the Li enhancements found here \citep{ba92}."520. Nevertheless. a significant number of J-stars (5%—105. LLovd-Evans 1991) show a very uniform 9.85 san emission which is believed to be due to the presence of a silicate dust shell (Little-Marenin.1936).," Nevertheless, a significant number of J-stars $5\%-10\%$, LLoyd-Evans 1991) show a very uniform 9.85 $\mu$ m emission which is believed to be due to the presence of a silicate dust shell \citep{lit86}."521. Also the detection of I5O masers in live J-ivpe stars has been reported (Engels1994)., Also the detection of $_2$ O masers in five J-type stars has been reported \citep{eng94}.522. This is rather strange. because silicate emission and ΠΟ masers," This is rather strange, because silicate emission and $_2$ O masers"523examining the intensity contours for the central region and taking the dvnamical center to be where tlie contours converged (o a maximum intensity.,examining the intensity contours for the central region and taking the dynamical center to be where the contours converged to a maximum intensity.524 Both the galaxy center and inclination angle were kept constant for all the rings., Both the galaxy center and inclination angle were kept constant for all the rings.525 Àn initial estimate of the position angle and the svstemic velocity of a galaxy was taken from NED. but it was clear [rom examination of the rotation curve fits and residuals that (he literature values needed to be changed for many ealaxies in our sample.," An initial estimate of the position angle and the systemic velocity of a galaxy was taken from NED, but it was clear from examination of the rotation curve fits and residuals that the literature values needed to be changed for many galaxies in our sample."526 We iterativelv changed the position angle aud (he systemic velocity ol a galaxy until the errors for the velocity [its in the individual rings in the rotation curve were a minimnunm., We iteratively changed the position angle and the systemic velocity of a galaxy until the errors for the velocity fits in the individual rings in the rotation curve were a minimum.527 The new values were all close to the literature values., The new values were all close to the literature values.528 The new improved parameters thal we adopted. are listed in Table 1., The new improved parameters that we adopted are listed in Table 1.529 As usual for the tilted-ring method. a rotation velocity is derived [ον each ring. weighting pixels with respect to the major axis of the ealaxyv.," As usual for the tilted-ring method, a rotation velocity is derived for each ring, weighting pixels with respect to the major axis of the galaxy."530 The width of a ring was set to (he mean synthesized beamwidth: for Nyquist sampling. we used rines spaced at half beam intervals.," The width of a ring was set to the mean synthesized beamwidth; for Nyquist sampling, we used rings spaced at half beam intervals."531" The rotational velocities at the bulge radii (/2,) and at the bar ends ως) were interpolated [rom the rotation curves.", The rotational velocities at the bulge radii $R_{blg}$ ) and at the bar ends $R_{bar}$ ) were interpolated from the rotation curves.532" We derive (he mass concentration as JCrap?Roar where (y, and (y, are the rotational velocities at the bulge radius and at the bar axis length respectively,"," We derive the mass concentration as ${v_{blg}}^{2}R_{blg}/{v_{bar}}^{2}R_{bar}$ , where $v_{blg}$ and $v_{bar}$ are the rotational velocities at the bulge radius and at the bar semi-major axis length respectively."533 This relation does not include (he effects of bulge ancl disk geometry which will be different [or each galaxy., This relation does not include the effects of bulge and disk geometry which will be different for each galaxy.534 It also assumes that the magnitude of elliplical streaming is not significant at the bulge radius or bar ends., It also assumes that the magnitude of elliptical streaming is not significant at the bulge radius or bar ends.535 This is à reasonable assumption in (he bulge where the potential is fairly axisvimnnetrie. and also at (he bar ends where the disk potential begins to be more important (han the bar.," This is a reasonable assumption in the bulge where the potential is fairly axisymmetric, and also at the bar ends where the disk potential begins to be more important than the bar."536 We discuss evidence supporting (his assuniption in (he next section., We discuss evidence supporting this assumption in the next section.537 The values of dderived for our sample of bars are shown in Table 2., The values of derived for our sample of bars are shown in Table 2.538 Also shown are the dvnamical masses within the bulge and bar of the galaxies., Also shown are the dynamical masses within the bulge and bar of the galaxies.539 Figure 1 shows the deprojected ellipticiüies in the plane of the galaxies. plotted against the central mass concentrationμοι.," Figure 1 shows the deprojected ellipticities in the plane of the galaxies, plotted against the central mass concentration."540 The errors lor both axes have been calculated using the standard error propagation equation. based on the uncertainties in the observed quantities (Devington Robinson 1992).," The errors for both axes have been calculated using the standard error propagation equation, based on the uncertainties in the observed quantities (Bevington Robinson 1992)."541 The error along the ellipticity axis includes a coefficient due to the deprojection of the bar onto the plane of the galaxy., The error along the ellipticity axis includes a coefficient due to the deprojection of the bar onto the plane of the galaxy.542 It is clear even [rom just visual inspection that there is a correlation between bar ellipticiGes ancl ii the galaxies., It is clear even from just visual inspection that there is a correlation between bar ellipticities and in the galaxies.543 However. for a more quantitative estimate of the correlation. we have determined a linear correlation coefficient [or the sample using (wo different methods.," However, for a more quantitative estimate of the correlation, we have determined a linear correlation coefficient for the sample using two different methods."544 An accurate estimate of the correlation coefficient should include the errors onboth axes.," An accurate estimate of the correlation coefficient should include the errors onboth axes,"545is followecl by an adiabatic deceleration. stage in. which the blast wave Lorentz factor decreases with radius A. as D.oxOU (Dlandford .Melxee. L976).,is followed by an adiabatic deceleration stage in which the blast wave Lorentz factor decreases with radius $\Rs$ as $\Gs \propto \Rs^{-3/2}$ (Blandford McKee 1976).546 The transition between the free expansion and deceleration phases occurs roughly at the racius Ay at which the energy in the swept-up surrounding material becomes comparable to £i: a where e is the total energy density of the surrounding material: fora typical interstellar medium of number density n. this is simply the rest-mass energy e8ηΠιο.," The transition between the free expansion and deceleration phases occurs roughly at the radius $\Rd$ at which the energy in the swept-up surrounding material becomes comparable to $\Egrb$: (, where $e$ is the total energy density of the surrounding material; for a typical interstellar medium of number density $n$, this is simply the rest-mass energy $e \approx n \, \mpr c^2$."547 The subsequent evolution of Ες is then given by and thus once 2.Ry the behaviour of LV. is independent of the value of 9.," The subsequent evolution of $\Gs$ is then given by ) ( , and thus once $\Rs > \Rd$ the behaviour of $\Gs$ is independent of the value of $\eta$."548 We now consider various scenarios for the acceleration of VHECRs by the relativistic blast wave of a fireball expanding into the general interstellar medium., We now consider various scenarios for the acceleration of UHECRs by the relativistic blast wave of a fireball expanding into the general interstellar medium.549 In order for particles to be accelerated beyond: the initial boost by shock acceleration. of the Fermi type. repeated shock crossings must occur.," In order for particles to be accelerated beyond the initial boost by shock acceleration of the Fermi type, repeated shock crossings must occur."550 The acceleration time. which we argued above is dominated. by the upstream: residence time /4. must therefore be shorter than the age of the blast wave measured in the upstream frame. 2.fe.," The acceleration time, which we argued above is dominated by the upstream residence time $\tu$, must therefore be shorter than the age of the blast wave measured in the upstream frame, $\Rs / c$."551 Using (8)). the maximum energy for which this can be the case is E XqDEU.- (C... ).," Using \ref{tup}) ), the maximum energy for which this can be the case is E q B (, ) ."552 The most. favourable regime cnereetically is that. of regular dellection upstream. which requires (i.c£2: in what follows we will assume that this is the case.," The most favourable regime energetically is that of regular deflection upstream, which requires $\lc \ga \Rs$; in what follows we will assume that this is the case."553 Note hat the maximum enerev (11)) is larger by a [actor lL. han that resulting from a simple ecometrical comparison of the evration radius with £2., Note that the maximum energy \ref{limit}) ) is larger by a factor $\Gs$ than that resulting from a simple geometrical comparison of the gyration radius with $\Rs$.554 This is due to the fact that a yarticle typically only executes a fraction ~Lt of a Larmor orbit upstream before recrossing the shock., This is due to the fact that a particle typically only executes a fraction $\sim\! \Gs^{-1}$ of a Larmor orbit upstream before recrossing the shock.555" The behaviour of E; as a function of 2. implies that the ughest energy in (11)) is reached at the transition between he free expansion ancl deceleration phases. £2,2 f?4."," The behaviour of $\Gs$ as a function of $\Rs$ implies that the highest energy in \ref{limit}) ) is reached at the transition between the free expansion and deceleration phases, $\Rs556\simeq \Rd$ ."557 Using he definition (9)) of R4. we obtain: Sh ZB«Sev. for ions of charge g=Ze. where Boy—Bf{10PG). S—Sotaf(107(aperg). n_SyEflO? and ng_—nf(1em/47).," Using the definition \ref{decradius}) ) of $\Rd$, we obtain: E 5 Z, for ions of charge $q = Z e$, where $B_{-6} \equiv B / \left(10^{-6}{\rm G}558\right)$, $\cE_{52} \equiv \Egrb / \left(10^{52} \,\erg\right)$, $\eta_{3}559\equiv \eta / 10^3$, and $n_0 \equiv n / \left(1 \,\cm^{-3}\right)$."560 Given tvpical interstellar magnetic fields of a few microgauss and the weak dependence on the other parameters. this rules out the production of CUILECRs by Fermi acceleration at the unmoclified external blast waves of relativistic fireballs.," Given typical interstellar magnetic fields of a few microgauss and the weak dependence on the other parameters, this rules out the production of UHECRs by Fermi acceleration at the unmodified external blast waves of relativistic fireballs."561 One remaining possibility for. blast wave acceleration. of ULLIECTs is that they result. from the initial shock crossing evcle energy boost of a pre-existing upstream population of relativistic particles., One remaining possibility for blast wave acceleration of UHECRs is that they result from the initial shock crossing cycle energy boost of a pre-existing upstream population of relativistic particles.562 This is. in essence. the first of the two acceleration mechanisms considered by. Vietri (1995).," This is, in essence, the first of the two acceleration mechanisms considered by Vietri (1995)."563" This initial boost. requires only the time £4- for the particle to be scattered: downstream. which if equipartition llds are generated. can be considerably shorter. than f,. yielding a corresponcinely higher. maximum. energy."," This initial boost requires only the time $\td$ for the particle to be scattered , which if equipartition fields are generated can be considerably shorter than $\tu$, yielding a correspondingly higher maximum energy."564 Requiring the downstream hall-evcle time (7)) for a particle of final energy. ££ to be shorter than the age of the fireball. we obtain ST ZSfLhj ov.," Requiring the downstream half-cycle time \ref{tdown}) ) for a particle of final energy $E$ to be shorter than the age of the fireball, we obtain E 7 Z ."565 This process can thus attain CLECR energies. provided particles with sullicient initial energy to be boosted in this range are present. upstream.," This process can thus attain UHECR energies, provided particles with sufficient initial energy to be boosted in this range are present upstream."566" To reach 107eV. this requires relativisticM particles. with. energy above ~107llo,4,246eV."," To reach $10^{20} \, \eV$, this requires relativistic particles with energy above $\sim \! 10^{14} \, \eta_3^{-2} \, \eV$."567 The energy invested in boosting pre-existing relativistic particles is of order L2ec per unit. volume swept up by the blast wave. where ecg is the upstream energy. density of relativistic particles.," The energy invested in boosting pre-existing relativistic particles is of order $\Gs^2 \ecr$ per unit volume swept up by the blast wave, where $\ecr$ is the upstream energy density of relativistic particles."568 Meanwhile. the blast wave expends an energv ~[7e per unit volume in shock-heating the surrounding medium. which has upstream energy density c8pnmo.," Meanwhile, the blast wave expends an energy $\sim \Gs^2 e$ per unit volume in shock-heating the surrounding medium, which has upstream energy density $e \approx n \mpr c^2$."569 Vhus the fraction of the fireball's energy. that can ego into boosting cosmic ravs is of order f—Confe., Thus the fraction of the fireball's energy that can go into boosting cosmic rays is of order $f \sim \ecr/e$.570" ‘Taking the interstellar medium. values in our Galaxy to be typical. Con~LeVem and ec10""nyeVem WO SCO that this mechanism can only have a very low clliciency. [£10""."," Taking the interstellar medium values in our Galaxy to be typical, $\ecr571\sim 1 \, \eV \, \cm^{-3}$ and $e \simeq 10^9 \, n_0 \, \eV \, \cm^{-3}$, we see that this mechanism can only have a very low efficiency, $f \la 10^{-9}$."572 The οποίοι vield. of ULIECIts. required in the GRB hypothesis. fz0.1 (Waxman 1995a: ιο 1905). could only be obtained if a lage fraction. of the surrounding mecdium's energv density was in relativistic particles.," The efficient yield of UHECRs required in the GRB hypothesis, $f \ga 0.1$ (Waxman 1995a; Vietri 1995), could only be obtained if a large fraction of the surrounding medium's energy density was in relativistic particles."573 This is in fact not implausible in the neutron star binary. merger scenario for GRBs. as we now argue.," This is in fact not implausible in the neutron star binary merger scenario for GRBs, as we now argue."574 The surrounding medium in which the relativistic fireball explodes is probably one modified. by the activity of the progenitor svstenir prior to the burst event., The surrounding medium in which the relativistic fireball explodes is probably one modified by the activity of the progenitor system prior to the burst event.575 In theneutronstar binary merger scenariofor CRBs. these progenitors are identified. with the binarypulsar svstems observed in our ownGalaxy (Naravan.Paczvásski Piran 1992).," In theneutronstar binary merger scenariofor GRBs, these progenitors are identified with the binarypulsar systems observed in our ownGalaxy (Narayan,Paczyńsski Piran 1992)."576 The pulsar in these systems can be expected to emit a, The pulsar in these systems can be expected to emit a577The situation changes dramatically once the cluster is endowed with a cistribution of stellar tnasses.,The situation changes dramatically once the cluster is endowed with a distribution of stellar masses.578 Here. binaries can form even where the deusity is close to the average.," Here, binaries can form even where the density is close to the average."579 We may demonstrate (his fact through a slight alteration of the heuristic derivation for /5 given in (2008)., We may demonstrate this fact through a slight alteration of the heuristic derivation for $t_b^\ast$ given in \citet{bt08}.580. Suppose that stars require a minimun 1inass 77 to be part of a binary., Suppose that stars require a minimum mass $m$ to be part of a binary.581 The time M for a eiven one of these objects to come withiu distance b of another with comparable mass is Here fj is the number [raction of such stars. 5n is the average cluster number deusity. aud σ the velocity dispersion.," The time $\Delta t$ for a given one of these objects to come within distance $b$ of another with comparable mass is Here $f_m$ is the number fraction of such stars, $n$ is the average cluster number density, and $\sigma$ the velocity dispersion."582 During this encounter. there is a probability that a third star will also be within the interaction distance 6.," During this encounter, there is a probability that a third star will also be within the interaction distance $b$."583 This star can have the average massΑΝ., This star can have the average mass.584.. Thus. the time for the original star to sulfer a binary-formiug triple encounter is aboutσ).," Thus, the time for the original star to suffer a binary-forming triple encounter is about."585". There are fy,N such stars iu the cluster."," There are $f_m\,N$ such stars in the cluster."586" The time foreng such star to form a binary is In order for a hard binary to form. the egravitational potential enereyOe of the binary must be equal to or greater than the average kinetic energy iu the cluster: Thus. From4 the virial⋅⋅ theorem.fry. where ry, is. the clusters. virial⋅⋅ radius."," The time for such star to form a binary is In order for a hard binary to form, the gravitational potential energy of the binary must be equal to or greater than the average kinetic energy in the cluster: Thus, From the virial theorem, where $r_v$ is the cluster's virial radius."587⋅ τονUsing this. expression along with the approximation that we [iud Now the relaxatiou time [rom equation (1) may be approximated as Dividing equation (10) by equation (11) yields, Using this expression along with the approximation that we find Now the relaxation time from equation (1) may be approximated as Dividing equation (10) by equation (11) yields588the stability can be tested.,the stability can be tested.589 In the case of instability either the treatment of the stability problem is insufficient or the system is in thermal imbalance., In the case of instability either the treatment of the stability problem is insufficient or the system is in thermal imbalance.590 In the case of stability the efficiency can be calibrated., In the case of stability the efficiency can be calibrated.591 For the system AB And (P?=0.531892 day) Hrivnak (1988)) determined the mass ratio y= 0.191. the mass ÀJ=1.5. and other properties listed 1n the first line of Table 4.. where temperatures are in. Kelvin and luminosities and radit are in solar units. (," For the system AB And $P=0.331892$ day) Hrivnak \cite{hri}) ) determined the mass ratio $q=0.491$ , the mass $M=1.5M_{\sun}$, and other properties listed in the first line of Table \ref{AB}, where temperatures are in Kelvin and luminosities and radii are in solar units. ("592This holds às well in the following tables.),This holds as well in the following tables.)593 The temperature is an average between the components., The temperature is an average between the components.594 According to Hrivnak the uncertainty in £ is large., According to Hrivnak the uncertainty in $L$ is large.595 Standard errors in the last digits are | for the radii and 3 for the degree of contact., Standard errors in the last digits are 1 for the radii and 3 for the degree of contact.596 Treating the observed values for period. mass. and mass ratio as constraints and using οι to adjust the mass ratio. a configuration is determined by the metallicity Z and the parameters fe. AT... Models for AB And are summarized in Table 4..," Treating the observed values for period, mass, and mass ratio as constraints and using $X_{{\rm c}1}$ to adjust the mass ratio, a configuration is determined by the metallicity $Z$ and the parameters $f_{\rm E},\Delta T_{\rm e}$ Models for AB And are summarized in Table \ref{AB}."597 The charge is positive and the models are indeed stable., The charge is positive and the models are indeed stable.598 This has been checked by evolutionary. calculations following a small perturbations., This has been checked by evolutionary calculations following a small perturbations.599 The amplitude of the circulation is small (f2 0.01)., The amplitude of the circulation is small $f\simeq 0.04$ ).600 The models show that the theory is compatible with the observations., The models show that the theory is compatible with the observations.601 The results for the temperature suggests the metallicity Z=0.011., The results for the temperature suggests the metallicity $Z=0.011$.602 The results forthe degree of contact suggest a very small value for the efficiency (fgc101)., The results forthe degree of contact suggest a very small value for the efficiency $f_{\rm E}\simeq 10^{-4}$ ).603 The model with these two values and with AT.=200 K is in excellent agreement with the observations. except for the radit which are somewhat too small.," The model with these two values and with $\Delta T_{\rm e}=200$ K is in excellent agreement with the observations, except for the radii which are somewhat too small."604 We shall return to thispoint., We shall return to thispoint.605 The geometry ofthe system in the equatorial plane is shown in Fig. 25.., The geometry of the system in the equatorial plane is shown in Fig. \ref{fig25}.606 Note that the fractional extent in radius of the convective envelope is much larger in the primary than in the secondary., Note that the fractional extent in radius of the convective envelope is much larger in the primary than in the secondary.607 The fractional extent in mass is also much larger., The fractional extent in mass is also much larger.608 We checked also the effects of changes inthe chemical profile., We checked also the effects of changes inthe chemical profile.609 In the models listed in Table 4. calculated with the standard profile Gr.=0.5). the hydrogen content V.4 is somewhat lower than 0.4.," In the models listed in Table \ref{AB}, calculated with the standard profile $x_{\rm c}=0.5$ ), the hydrogen content $X_{{\rm c}1}$ is somewhat lower than 0.4."610 Models with ου=0.3 give similar results with Vy0.2., Models with $x_{\rm c}=0.3$ give similar results with $X_{{\rm c}1}\simeq 0.2$.611 Accordingly. the choice of the chemical profile is of minor importance.," Accordingly, the choice of the chemical profile is of minor importance."612 The amount of hydrogen already converted to helium is also similar and can be used to estimate the age of the system., The amount of hydrogen already converted to helium is also similar and can be used to estimate the age of the system.613 The contact binaries BV. Dra and BW Dra are particularly interesting since they form à. visual. binary., The contact binaries BV Dra and BW Dra are particularly interesting since they form a visual binary.614 Precise observational results have been obtained by Kaluzny Rueinski (1986))., Precise observational results have been obtained by Kaluzny Rucinski \cite{karu}) ).615 We begin with BV Dra., We begin with BV Dra.616 Period (0.350067 day). mass (1.TAL: ) and mass ratio (y.= 0.111) are again treated as constraints.," Period (0.350067 day), mass $1.47M_{\sun}$ ) and mass ratio $q=0.411$ ) are again treated as constraints."617 Other observational results are listed in the first line of Table 5.., Other observational results are listed in the first line of Table \ref{BV}.618 Standard errors in the last digits are 2] for the luminosity. | for the radii. and 27 for the degree of contact.," Standard errors in the last digits are 21 for the luminosity, 1 for the radii, and 27 for the degree of contact."619 The models of BV turned again out to be stable., The models of BV turned again out to be stable.620 Comparing the temperatures with the observed value we obtain the metallicity Z=0.006. in accordance with the result of Kaluzny Rucinski.," Comparing the temperatures with the observed value we obtain the metallicity $Z=0.006$, in accordance with the result of Kaluzny Rucinski."621 The degree of contact suggests again a very low efficiency (fpc10 °)., The degree of contact suggests again a very low efficiency $f_{\rm E}\simeq 10^{-3}$ ).622 The model with these values and with A7;.=200 K is in very good agreement with the observations. again apart from the fact that the radii are somewhat too small.," The model with these values and with $\Delta T_{\rm e}=200$ K is in very good agreement with the observations, again apart from the fact that the radii are somewhat too small."623 The constraints for thissystem are P?=0.292167 day. AL= :.4= 0.280.," The constraints for thissystem are $P=0.292167$ day, $M=1.18 M_{\sun}$ , $q=0.280$ ."624 Other observational results are listed in the first line of Table 6.. Standard errors in the last digits are, Other observational results are listed in the first line of Table \ref{BW}.. Standard errors in the last digits are625were detected inIi.,were detected in.626 Morganti et al., Morganti et al.627 find that ccan be present in different morphologies: deises seem to be as common as oll-set clouds ancl tails. though they occur mostly in the relatively gas-rich svstenis.," find that can be present in different morphologies: discs seem to be as common as off-set clouds and tails, though they occur mostly in the relatively gas-rich systems."628 Recently rotation curves of clelises in low surface brightness galaxies and cwarf galaxies. complemented. with Lla observations. have been used. not only to confirm the existence of dark matter haloes. but also to obtain estimates on the inner slope of the density profiles of the haloes (c.g. van den Bosch ct al. 2000:," Recently rotation curves of discs in low surface brightness galaxies and dwarf galaxies, complemented with $\alpha$ observations, have been used not only to confirm the existence of dark matter haloes, but also to obtain estimates on the inner slope of the density profiles of the haloes (e.g. van den Bosch et al. \nocite{2000AJ....119.1579V};"629 Welclrake. cle Blok Walter 2003)).," Weldrake, de Blok Walter \nocite{2003MNRAS.340...12W}) )."630 Simulations within a cold clark matter (CDM) cosmology viel haloes with cusps in their centres (NEW profiles. see. Navarro. Freonk White 1996)). but observations suggest core-clominatecl profiles (e.g. ce Blok Bosma 2002:: de Blok 2005)).," Simulations within a cold dark matter (CDM) cosmology yield haloes with cusps in their centres (NFW profiles, see Navarro, Frenk White \nocite{1996ApJ...462..563N}) ), but observations suggest core-dominated profiles (e.g. de Blok Bosma \nocite{2002A&A...385..816D}; de Blok \nocite{2005ApJ...634..227D}) )."631 Detailed studies of rotation curves of earlv-tvpe galaxies that contain dedises are sparser. due to lack of spatial resolution: to detect low ssurface densities. larger beams are needed.," Detailed studies of rotation curves of early-type galaxies that contain discs are sparser, due to lack of spatial resolution: to detect low surface densities, larger beams are needed."632 Also. only [ew earlv-tvpe galaxies have deises that are extended: and regular enough to allow for detailed: studies.," Also, only few early-type galaxies have discs that are extended and regular enough to allow for detailed studies."633 Comparing ML values at. large. radii derived from vvelocities. to AAL at smaller radii measured. [rom ionised eas kinematics. the conclusion is that earlv-tvpe. galaxies also have dark matter dominated haloes (e.g. Bertola ct al. 1993:," Comparing $M/L$ values at large radii, derived from velocities, to $M/L$ at smaller radii measured from ionised gas kinematics, the conclusion is that early-type galaxies also have dark matter dominated haloes (e.g. Bertola et al. \nocite{1993ApJ...416L..45B};"634 Morganti et al. 1997:: , Morganti et al. \nocite{ 1997AJ....113..937M}; ;635Sadler et al. )00)., Sadler et al. \nocite{2000AJ....119.1180S}) ).636 Franx. van Gorkom de Zeeuw used the rring of the elliptical galaxy. I€ 2006 to determine not only the mass. but also the shape of the dark halo.," Franx, van Gorkom de Zeeuw \nocite{1994ApJ...436..642F} used the ring of the elliptical galaxy IC 2006 to determine not only the mass, but also the shape of the dark halo."637 They concluced that LC 2006 is surrounded by an axisvmmoetric dark halo. using the geometry of the ring and an harmonic expansion of its velocity map.," They concluded that IC 2006 is surrounded by an axisymmetric dark halo, using the geometry of the ring and an harmonic expansion of its velocity map."638 In this paper. we present a similar analysis of the regularly rotating rring around the elliptical (124) field galaxy NGC 2974.," In this paper, we present a similar analysis of the regularly rotating ring around the elliptical (E4) field galaxy NGC 2974."639 Kim et al., Kim et al.640 observed this galaxy before in but their data hack lower spatial resolution than ours. and they found. a filled. clise instead. of a ring.," \nocite{1988ApJ...330..684K} observed this galaxy before in but their data had lower spatial resolution than ours, and they found a filled disc instead of a ring."641 Cinzano van der Marel found an embedded: stellar disc in their cdvnamical model of this galaxy. based. upon long-slit spectroscopic data. but. Emsellem.. Goudfrooij Ferrutt constructed a dvnamical model of NGC 2974 based on TIGE integral-field spectrography and. long-slit stellar kinematics. that does not require a hidden disc structure.," Cinzano van der Marel \nocite{{1994MNRAS.270..325C}} found an embedded stellar disc in their dynamical model of this galaxy, based upon long-slit spectroscopic data, but Emsellem, Goudfrooij Ferruit \nocite{2003MNRAS.345.1297E} constructed a dynamical model of NGC 2974 based on TIGER integral-field spectrography and long-slit stellar kinematics, that does not require a hidden disc structure."642 They did report the detection of a two-arm gaseous spiral in the inner 200 pe of NGC 2974 from high resolution WEPC2 imaging., They did report the detection of a two-arm gaseous spiral in the inner 200 pc of NGC 2974 from high resolution WFPC2 imaging.643 WKrajnovié et al., Krajnović et al.644 constructed. axisvmmoetric dynamical models of both the stars ane ionisecl gas based upon integral-field. cata., \nocite{2005MNRAS.357.1113K} constructed axisymmetric dynamical models of both the stars and ionised gas based upon integral-field data.645" ""These models require a component with high angular momentum. consisting of à somewhat Hattened: distribution of stars. though not a thin stellar disc."," These models require a component with high angular momentum, consisting of a somewhat flattened distribution of stars, though not a thin stellar disc."646 LEmisellem et al., Emsellem et al.647" classify NGC. 2974 as a Last rotator, which means that it possesses large-scale rotation and that its angular momentum is well defined."," \nocite{2007MNRAS.379..401E} classify NGC 2974 as a fast rotator, which means that it possesses large-scale rotation and that its angular momentum is well defined."648 Some of the characteristics of NGC 2974 are given in Table 1.., Some of the characteristics of NGC 2974 are given in Table \ref{tab:ngc2974}.649 For our analysis of NCC 2974 we combine kinematics of neutral gas. obtained [rom our observations with the Very Large Array (VLA). with that of ionised gas. obtained with the integral-field spectrograph ((Bacon οἱ al. 2001)).," For our analysis of NGC 2974 we combine kinematics of neutral gas, obtained from our observations with the Very Large Array (VLA), with that of ionised gas, obtained with the integral-field spectrograph (Bacon et al. \nocite{2001MNRAS.326...23B}) )."650 This combination of a small scale two-dimensional gas velocity map in the centre of the galaxy. anc alli vvelocity map at the outskirts. allows measurements of a rotation curve ranging from LOO pe within the centre of the galaxy to 10 kpe at the edges of the rring.," This combination of a small scale two-dimensional gas velocity map in the centre of the galaxy, and a velocity map at the outskirts, allows measurements of a rotation curve ranging from 100 pc within the centre of the galaxy to 10 kpc at the edges of the ring."651 We use this rotation curve. together with erouncl- and space based optical imaging. to determine the dark matter content in NGC 2974. and to constrain the shape of the dark halo.," We use this rotation curve, together with ground- and space based optical imaging, to determine the dark matter content in NGC 2974, and to constrain the shape of the dark halo."652 In section 2. we discuss the two datasets and. their reduction. and describe the rring.," In section 2, we discuss the two datasets and their reduction, and describe the ring."653 We concentrate on the analysis of the velocity maps in section 3., We concentrate on the analysis of the velocity maps in section 3.654 Section 4 is devoted to the rotation curve that we extract from the velocity maps. and in section 5 we show mass mocdels with various halo mocoels. and find the best fit to the rotation curve.," Section 4 is devoted to the rotation curve that we extract from the velocity maps, and in section 5 we show mass models with various halo models, and find the best fit to the rotation curve."655 Section 6 summarizes our results., Section 6 summarizes our results.656 Earher VLA observations (Ixim ct al. 1988)), Earlier VLA observations (Kim et al. \nocite{1988ApJ...330..684K}) )657 of NGC 2974 showed that this galaxy contains a significant amount of tthat. in their observations. appears to be distributed. in a regularly rotating disc.," of NGC 2974 showed that this galaxy contains a significant amount of that, in their observations, appears to be distributed in a regularly rotating disc."658 Given the modest spatial ancl velocity resolution of those observations. we re-observed NGC 2974 with the VLA C-array while also using a cilferent frequency setup that allows us to study this galaxy at both higher spatial and higher velocity. resolution.," Given the modest spatial and velocity resolution of those observations, we re-observed NGC 2974 with the VLA C-array while also using a different frequency setup that allows us to study this galaxy at both higher spatial and higher velocity resolution."659 The observations were performed on 11 and 19 September 2005 with a total on-source integration time of 15 hours., The observations were performed on 11 and 19 September 2005 with a total on-source integration time of 15 hours.660 In each observation. two partially overlapping bands of 3.15 MllIz and 64 channels were used.," In each observation, two partially overlapping bands of 3.15 MHz and 64 channels were used."661 Phe two bands were ollset. by 500 iin central velocity., The two bands were offset by 500 in central velocity.662 This frequency setup allows us to obtain eood velocity resolution over a wide range of velocities (about 1080 1)., This frequency setup allows us to obtain good velocity resolution over a wide range of velocities (about 1080 ).663 The cata were calibrated following standard procedures using the ΛΗΛ software package (Sault. TeubenWright 19953).," The data were calibrated following standard procedures using the MIRIAD software package (Sault, TeubenWright \nocite{xxx}) )."664 A spectral-line data cube was mace using, A spectral-line data cube was made using665changes little in the time taken for the photon to travel through it. ó averages to 0.,"changes little in the time taken for the photon to travel through it, $\delta$ averages to $0$."666 Thus we impose the condition where A is any parameterization along the geodesic., Thus we impose the condition where $\lambda$ is any parameterization along the geodesic.667 Finally. h(r.t)=O at both the source and the observer's locations since the only effect under consideration here is due to matter inhomogeneities between the source and the observer.," Finally, $h(r,t)=0$ at both the source and the observer's locations since the only effect under consideration here is due to matter inhomogeneities between the source and the observer."668" Following Peebles(1993.page276) (but note the different definition of // which allows the calculation to be carried out to all orders). the stress energy conservation law leads to or Since pi/p,2—3d/a in the unperturbed universe. which has solution The integration constant has been determined by requiring that when 620. //z0."," Following \citet[page 276]{peebles93} (but note the different definition of $h$ which allows the calculation to be carried out to all orders), the stress energy conservation law leads to or Since $\dot{\rho}_b / \rho_b = -3\dot{a} / a$ in the unperturbed universe, which has solution The integration constant has been determined by requiring that when $\delta=0$, $h=0$."669 By integrating along a radial null geodesic from emission at (f.r7)=(4.0) to observation at (6.7)=Gas.7). the radial coordinate where a photon arrives at the observer in the perturbed universe can be determined and compared to the radial coordinate of the observer in the FRW universe.," By integrating along a radial null geodesic from emission at $(t,r)=(t_{\mathrm{em}},0)$ to observation at $(t,r)=(t_{\mathrm{obs}},r)$, the radial coordinate where a photon arrives at the observer in the perturbed universe can be determined and compared to the radial coordinate of the observer in the FRW universe."670 In the perturbed FRW universe Equation (10)) gives a differential equation which may be solved for rf)., In the perturbed FRW universe Equation \ref{eq_pert_de}) ) gives a differential equation which may be solved for $r(t)$ .671 The position of the photon is completely described by the function. (tf)., The position of the photon is completely described by the function $r(t)$.672 A(t) is now defined along the geodesic as A(t)=AG(r).rp.," $h(t)$ is now defined along the geodesic as $h(t)\equiv h(r(t),t)$."673 Similarly. 6(4)=1). 1).," Similarly, $\delta(t)\equiv \delta(r(t),t)$ ."674 Rearranging equation (10)) and integrating. one obtains Similarly. 1n the FRW universe Substituting for 1—/7 from equation (9)) and making the second order approximation (1-6)?=146/3—67/9. equation (11) becomes Equations (12)) and (13)) are both written using comoving coordinates and proper time and so may be compared directly.," Rearranging equation \ref{eq_pert_de}) ) and integrating, one obtains Similarly, in the FRW universe Substituting for $1-h$ from equation \ref{eq_def_h}) ) and making the second order approximation $(1+\delta)^{1/3} \approx6751+\delta/3-\delta^2/9$, equation \ref{eq_pert_geodesic_1}) ) becomes Equations \ref{eq_FRW_geodesic}) ) and \ref{eq_pert_geodesic_2}) ) are both written using comoving coordinates and proper time and so may be compared directly."676 The first term in equation (13)) may be replaced using equation (12)) and the second term is 0 due to equation (5)) leaving Since 6->0 and «>0 then rpgw2ron. and rus=nae only if à=0 at all points along the lightcone.," The first term in equation \ref{eq_pert_geodesic_2}) ) may be replaced using equation \ref{eq_FRW_geodesic}) ) and the second term is $0$ due to equation \ref{eq_average}) ) leaving Since $\delta^2 \ge 0$ and $a > 0$ then $r_{\mathrm{FRW}} \ge r_{\mathrm{obs}}$ and $r_{\mathrm{obs}} =677r_{\mathrm{FRW}}$ only if $\delta=0$ at all points along the lightcone."678 We may use this to compare the observed number of photons 1n the perturbed universe and the FRW universe., We may use this to compare the observed number of photons in the perturbed universe and the FRW universe.679 From equation (1)). Telescopes in a perturbed FRW universe receive on average more photons from a source at a given redshift than telescopes with the same area in a FRW universe and therefore have a higher apparent magnitude.," From equation \ref{eq_photon_number}) ), Telescopes in a perturbed FRW universe receive on average more photons from a source at a given redshift than telescopes with the same area in a FRW universe and therefore have a higher apparent magnitude."680 This is the main result of this paper., This is the main result of this paper.681 The focusing theorem (Schneideretal.1992. shows that a light beam is magnified if it is affected by gravitational lensing but does not go through a caustic., The focusing theorem \citep[page 132]{SEF92} shows that a light beam is magnified if it is affected by gravitational lensing but does not go through a caustic.682 In light of the focusing theorem. the result in Section 3. is not surprising.," In light of the focusing theorem, the result in Section \ref{sec_z} is not surprising."683 The conclusion of Section 3. is significant because it shows that gravitational lensing can cause magnification for all observers without violating conservation of photon number., The conclusion of Section \ref{sec_z} is significant because it shows that gravitational lensing can cause magnification for all observers without violating conservation of photon number.684 The calculation in Section 3 is very general in that it does not depend up the distribution of matter. only that there are matter perturbations.," The calculation in Section \ref{sec_z} is very general in that it does not depend up the distribution of matter, only that there are matter perturbations."685 It does assume that the global structure of the universe is FRW and that the departure from FRW is slight., It does assume that the global structure of the universe is FRW and that the departure from FRW is slight.686 It seems plausible that any greater departure from FRW will not remove the effect., It seems plausible that any greater departure from FRW will not remove the effect.687 Furthermore. the argument is based on the non-linearity of the relationship between matter and the metric. so It is easy to see how it may appliedin models other than FRW.," Furthermore, the argument is based on the non-linearity of the relationship between matter and the metric, so it is easy to see how it may appliedin models other than FRW."688 | am grateful to L. Ryder. J. Adams.W. Joyce. S.Seunarine and S. Bester for discussions and careful reading of drafts.," I am grateful to L. Ryder, J. Adams,W. Joyce, S.Seunarine and S. Besier for discussions and careful reading of drafts."689Fig.,Fig.690 9 shows some poloidal magnetic field lines and the radial velocity in the 2.5D and 3D jet., \ref{fig:mflines} shows some poloidal magnetic field lines and the radial velocity in the 2.5D and 3D jet.691" The latter was ""axisymmetrized"" by averaging over the azimuthal coordinate φ.", The latter was “axisymmetrized” by averaging over the azimuthal coordinate $\varphi$.692" In the outer, high-9, part of the 2.5D jet, where the toroidal field is especially strong and jet acceleration most efficient, there is an increase of the angular separation between the field lines which is absent in the 3D jet."," In the outer, $\vartheta$, part of the 2.5D jet, where the toroidal field is especially strong and jet acceleration most efficient, there is an increase of the angular separation between the field lines which is absent in the 3D jet."693" In other words, the poloidal magnetic flux decreases locally faster with distance in the 2.5D case."," In other words, the poloidal magnetic flux decreases locally faster with distance in the 2.5D case."694" To identify the accelerating forces that generate the kinetic energy flow discussed in Sect. 3.2,,"," To identify the accelerating forces that generate the kinetic energy flow discussed in Sect. \ref{sec:energy},"695" we compute P=fF. VdV, the instantaneous power (rate of work) delivered by a specific force F in the direction of the flow in the integrated volume."," we compute $P = \int \vec{F} \cdot \vec{v}696\, \de V$ , the instantaneous power (rate of work) delivered by a specific force $\vec{F}$ in the direction of the flow in the integrated volume."697" The combination of gas pressure and gravitational forces, -Vp−accounts for about one third of the power delivered by the sum of all forces in the whole volume, the corresponding acceleration takes place mainly below r ~30(sonic surface)."," The combination of gas pressure and gravitational forces, $-\Nabla p698-\rho\Nabla\Phi$, accounts for about one third of the power delivered by the sum of all forces in the whole volume, the corresponding acceleration takes place mainly below $r \approx 30$ (sonic surface)."699" The rest is accounted for by the Lorentz force, which we decompose as follows: with B,-Bye, and B,=B-B,."," The rest is accounted for by the Lorentz force, which we decompose as follows: with $\vec{B}_\varphi = B_\varphi \evarphi$ and $\vec{B}_\poloidal = \vec{B} -700\vec{B}_\varphi$."701 The corresponding components of P are plotted in Fig., The corresponding components of $P$ are plotted in Fig.702 10 as a function of the upper integral limit., \ref{fig:powers} as a function of the upper integral limit.703 The last force in Eq., The last force in Eq.704 has only an azimuthal component., has only an azimuthal component.705" It is important mainly below the radius (r~ 60), exerting a torque in the same direction in which rotation is applied at the lower boundary."," It is important mainly below the radius $r \sim 60$ ), exerting a torque in the same direction in which rotation is applied at the lower boundary."706 The next-to-last force vanishes in the axisymmetric case., The next-to-last force vanishes in the axisymmetric case.707" In the general case, it has only non-azimuthal components."," In the general case, it has only non-azimuthal components."708" Unlike the last force, it works against the flow; the two forces largely cancel each other in the 3D simulation."," Unlike the last force, it works against the flow; the two forces largely cancel each other in the 3D simulation."709" The Lorentz force associated with B, [second term", The Lorentz force associated with $B_\varphi$ [second term710"The triplet comprises three spectral lines which. according to the (NIST). are the ""blue? component and the ""red"" components at aand lines.","The triplet comprises three spectral lines which, according to the (NIST), are the “blue” component and the “red” components at and lines."711 The latter two are blended and consequently appear as a single spectral line 1n the solar spectrum., The latter two are blended and consequently appear as a single spectral line in the solar spectrum.712 The formation height of the Helium triplet happens in the upper chromosphere (?) and therefore it is especially interesting for the study of filaments and their magnetic properties. as already proved by many authors (????2?)..," The formation height of the Helium triplet happens in the upper chromosphere \citep{avrett94} and therefore it is especially interesting for the study of filaments and their magnetic properties, as already proved by many authors \citep{lin98,trujillo02,merenda07,casini09,kuckein09,sasso11}."713 The strong photospheric absorption line at ooriginates between the terms P;—?P» and has a Landé factor of gar=1.5.," The strong photospheric absorption line at originates between the terms $^3\mathrm{P}_2 \rightarrow\, ^3\mathrm{P}_2$ and has a Landé factor of $g_\mathrm{eff} = 1.5$."714 The combination of these lines is an excellent diagnostic tool to study magnetic fields and their coupling between the photosphere and the chromosphere., The combination of these lines is an excellent diagnostic tool to study magnetic fields and their coupling between the photosphere and the chromosphere.715 In this paper we also report on SOHO//MDI images which are used to understand the long term (days) evolution of the AR in the PIL region., In this paper we also report on /MDI images which are used to understand the long term (days) evolution of the AR in the PIL region.716 Active region NOAA 10781 emerged some weeks before our observing run oi the back side of the Sun., Active region NOAA 10781 emerged some weeks before our observing run on the back side of the Sun.717 Its magnetic configuration. às seen by SOHO//MDI (see below). clearly corresponds to that of an AR that is in its decay phase. 1.8. a round leader sunspot followed by facular regions of both polarities that show the latitudinal shear produced by the action of surface flows (e.g..?)..," Its magnetic configuration, as seen by /MDI (see below), clearly corresponds to that of an AR that is in its decay phase, i.e., a round leader sunspot followed by facular regions of both polarities that show the latitudinal shear produced by the action of surface flows \citep[e.g.,][]{vanballe08}."718 The filament is found above the PIL in the plage region., The filament is found above the PIL in the plage region.719 The time sequence of continuum images of MDI between 2005 July 2nd and 6th is presented in Fig. 2.., The time sequence of continuum images of MDI between 2005 July 2nd and 6th is presented in Fig. \ref{Fig:MDIcont}.720 Since filaments are not visible here. we expanded the FOV in order to use the AR leader sunspot. located in the lower right corner. as a reference.," Since filaments are not visible here, we expanded the FOV in order to use the AR leader sunspot, located in the lower right corner, as a reference."721 As shown in Fig. 2..," As shown in Fig. \ref{Fig:MDIcont},"722 it is not until July 4th that small pores start to emerge and gather together., it is not until July 4th that small pores start to emerge and gather together.723 White light images of the (TRACE)) confirm this behavior., White light images of the ) confirm this behavior.724 On July 5th. we increased the cadence in Fig.," On July 5th, we increased the cadence in Fig."725 2 to show that more pores quickly emerged and orphan penumbrae formed (see black rectangle on the July Sth 06:24 UT panel).," \ref{Fig:MDIcont}726 to show that more pores quickly emerged and orphan penumbrae formed (see black rectangle on the July 5th 06:24 UT panel)."727 Interestingly. these orphan penumbral regions act like a bridge connecting different small groups of pores together.," Interestingly, these orphan penumbral regions act like a bridge connecting different small groups of pores together."728 However. during July 6th the pores and orphan penumbrae disappear completely.," However, during July 6th the pores and orphan penumbrae disappear completely."729 It is also worth noting that the leading sunspot of the AR. which on July 2nd had a round and rather symmetric shape. also decays away slowly over this period of time and almost vanishes by the end of July 6th.," It is also worth noting that the leading sunspot of the AR, which on July 2nd had a round and rather symmetric shape, also decays away slowly over this period of time and almost vanishes by the end of July 6th."730 Figure 3 provides line-of-sight (LOS) magnetograms from MDI starting on 2005 July Ist until July 7th., Figure \ref{Fig:MDIevol} provides line-of-sight (LOS) magnetograms from MDI starting on 2005 July 1st until July 7th.731 This period was chosen in order to carry out a detailed study of the magnetic morphology and evolution of the AR., This period was chosen in order to carry out a detailed study of the magnetic morphology and evolution of the AR.732 Note that the FOV is smaller and the panels show different dates and times than Fig. 2.., Note that the FOV is smaller and the panels show different dates and times than Fig. \ref{Fig:MDIcont}.733 The AR and its PIL are well defined., The AR and its PIL are well defined.734 The image shows the two polarity regions as of July Ist., The image shows the two polarity regions as of July 1st.735" On the second day. the AR became more compact and the gray area in between the two opposite polarities. 1.8. the PIL. developed a winding shape that pointed approximately in the N-S direction,"," On the second day, the AR became more compact and the gray area in between the two opposite polarities, i.e. the PIL, developed a winding shape that pointed approximately in the N-S direction."736 The next two images (from July 3rd) show the AR becoming even more compact. since the black and white polarities approached each other.," The next two images (from July 3rd) show the AR becoming even more compact, since the black and white polarities approached each other."737 It is important to stress that. on this day. the AR rotated counterclockwise and the PIL oriented itself at ~45° with respect to solar north.," It is important to stress that, on this day, the AR rotated counterclockwise and the PIL oriented itself at $\sim 45^\circ$ with respect to solar north."738 These images also reveal that, These images also reveal that739"mag in A, and in slightly redder clusters. (JAK.)= nig as compared to (J0AK.)=0.8 mag for the general population of clusters.","mag in $K_s$ and in slightly redder clusters, $(J-K_s)=1.1$ mag as compared to $(J-K_s)=0.8$ mag for the general population of clusters."740 This confirms our results presented in Clarketal.(2007).. in which we performed a similar study with the smaller sample of LI. counterparts.," This confirms our results presented in \citet{cla07}, in which we performed a similar study with the smaller sample of IR counterparts."741 Following the work in Clarketal.(2008).. we explored the relationship between cluster mass ancl the detected number of X-ray sources.," Following the work in \citet{cla08}, we explored the relationship between cluster mass and the detected number of X-ray sources."742 In Clarketal.(2008) we defined a function 4 relating the number of X-ray detections per mass as a function of cluster mass., In \citet{cla08} we defined a function $\eta$ relating the number of X-ray detections per mass as a function of cluster mass.743 We found η is consistent with a mean value of Ηντος10., We found $\eta$ is consistent with a mean value of $\eta(F_{K_s})=5.5\times10^{-8}$.744 Using a A7 test. we compared a fitted slope of the plotted CF.) values to the mean a(n) value.," Using a $\chi^2$ test, we compared a fitted slope of the plotted $\eta(F_{K_s})$ values to the mean $\eta(F_{K_s})$ value."745 We found a NX?=0.45., We found a $\Delta\Sigma\chi^2=0.45$.746 Considering we found a X4?=19 in Clarketal.(2008).. our new study shows a stronger relation between g(£y.) ancl the mean value. ollectively ruling out any inclination towarcls more X-ray sources residing in more massive clusters other than through simple scaling arguments.," Considering we found a $\Delta\Sigma\chi^2=1.9$ in \citet{cla08}, our new study shows a stronger relation between $\eta(F_{K_s})$ and the mean value, effectively ruling out any inclination towards more X-ray sources residing in more massive clusters other than through simple scaling arguments."747" Inclucing the 27 star cluster counterparts to X-ray sources seen across all CDVLJRN, bands. we fit Muzual-Charlot (BC:Bruzual&Charlot2003). speetrophotonietric moclels to 10 0 “these clusters."," Including the 27 star cluster counterparts to X-ray sources seen across all $UBVIJK_s$ bands, we fit Bruzual-Charlot \citep[BC;][]{bru03} spectrophotometric models to 10 of these clusters."748" The DX “model fits indicate he X-rav-source-associated clusters are 710' 2. AL, . inmass.~9.LO"" 5210 vrin age. with extinction I0""varving )otweensd=0.9 1.5 mag."," The BC model fits indicate the X-ray-source-associated clusters are $7\times10^5$ – $2\times10^6$ $_{\sun}$ in mass, $\sim9\times10^6$ – $2\times10^7$ yr in age, with extinction varying between $A_V=0.3$ – 1.5 mag."749 These proprieties indicate star cluster counterparts to X-ray sources in the Antennae tend o be voung and massive. which are consistent with these X-ray sources being associated with star formation.," These proprieties indicate star cluster counterparts to X-ray sources in the Antennae tend to be young and massive, which are consistent with these X-ray sources being associated with star formation."750 In Ptaketal. (2006)... these authors also found similar mass results or their counterparts to ULXs.," In \citet{pta06}, these authors also found similar mass results for their counterparts to ULXs."751 While we can use multiwavelength photometry to describe cluster properties. there remains some uncertainty in these characteristics due to errors in magnitude ancl mocel limitations.," While we can use multiwavelength photometry to describe cluster properties, there remains some uncertainty in these characteristics due to errors in magnitude and model limitations."752 In. future work. we plan to acquire spectra of Antennae cluster counterparts ancl refine estimations of their properties.," In future work, we plan to acquire spectra of Antennae cluster counterparts and refine estimations of their properties."753 Vhe authors thank the stalfol Palomar Observatory for their excellent. assistance in Commissioning WIRC and obtaining these data., The authors thank the staff of Palomar Observatory for their excellent assistance in commissioning WIRC and obtaining these data.754 WIBC was mace possible by support from the NSE (NSE-ASTO32328522)) the Norris Foundation. and Cornell University.," WIRC was made possible by support from the NSF (NSF-AST0328522), the Norris Foundation, and Cornell University."755 Dased. on observations made with the NASA/ESA Llubble Space Telescope. obtained. from. the data archive at the. Space Telescope Science. Institute.," Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute."756 SIEScb is operated. by. the Association of Universities. [or th in Astronomy. Ine. under NASA contract NAS 5-26555.," STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555."757s DAIC is grateful for the many useful. discussions wit Michelle Eclwards and Valerie Mikles., DMC is grateful for the many useful discussions with Michelle Edwards and Valerie Mikles.758 SSE ancl DAIC received support in part by an NSP CAREER award (NSE- and an NSP grant. (NSE-ASTO507547)., SSE and DMC received support in part by an NSF CAREER award (NSF-9983830) and an NSF grant (NSF-AST0507547).759 We also thank J. Llouck for his support of the WIRC instrument project., We also thank J. Houck for his support of the WIRC instrument project.760the following and luminosity relations in dependence of the spectral subtype.,the following and luminosity relations in dependence of the spectral subtype.761 Like in the other two cases the spectral type grid is included in Table | and it is shown in of Fig. 1.., Like in the other two cases the spectral type grid is included in Table \ref{tab:LTgrid} and it is shown in of Fig. \ref{fig:sptypes}.762 As is visible in Fig., As is visible in Fig.763 | and Table |.. the SMC grid is shifted to temperatures compared to the LMC grid.," \ref{fig:sptypes} and Table \ref{tab:LTgrid}, the SMC grid is shifted to temperatures compared to the LMC grid."764 The reason for this is non-trivial. and deserves a thorough analysis using SMC metallicity NLTE atmospheres. which is beyond the scope of this paper.," The reason for this is non-trivial, and deserves a thorough analysis using SMC metallicity NLTE atmospheres, which is beyond the scope of this paper."765 Note that the sample only includes one star with a spectral type earlier than Os., Note that the sample only includes one star with a spectral type earlier than O4.766 This star is not included in the relations. but the relations are used from O2 to O9.5.," This star is not included in the relations, but the relations are used from O2 to O9.5."767 In order to assign a certain spectral subtype and luminosity class to a specific evolutionary point in time. stellar evolution models from (see also appendix Appendix A:)) are followed throughout the Z.-L-diagram.," In order to assign a certain spectral subtype and luminosity class to a specific evolutionary point in time, stellar evolution models from (see also appendix \ref{app:evol}) ) are followed throughout the $T_\mathrm{eff}$ $L$ -diagram."768 This is visualized in Fig. 2..," This is visualized in Fig. \ref{fig:tracks},"769 where the luminosity- and Z;j-evolution for six rotating as well as non-rotating stellar models from 20 to 120 M.. are plotted (rotating models as and non-rotating ones as lines)., where the luminosity- and $T_\mathrm{eff}$ -evolution for six rotating as well as non-rotating stellar models from 20 to 120 $M_\odot$ are plotted (rotating models as and non-rotating ones as ).770 In this figure. only the part of the evolution before the Wolf-Rayet (WR) stage is depicted.," In this figure, only the part of the evolution before the Wolf-Rayet (WR) stage is depicted."771 Objects that are still assumed to be core hydrogen are classified as WNL stars(?)., Objects that are still assumed to be core hydrogen are classified as WNL stars.772. However. wenote that the WNL classification is only used in appendix AppendixB:..," However, wenote that the WNL classification is only used in appendix \ref{app:stevol}."773 The present-day mass of a model during its evolutio through a spectral subelass is from now on referred to as evolutionary mass (Mat)., The present-day mass of a model during its evolution through a spectral subclass is from now on referred to as evolutionary mass $M_\mathrm{evol}$ ).774 Table 4 shows the new spectral type mass conversion. based on solar metallicity rotating evolutionary models from 10 to 120 M.(?)..," Table \ref{tab:Orot} shows the new spectral type mass conversion, based on solar metallicity rotating evolutionary models from 10 to 120 $M_\odot$."775" The rotating models have initial rotational velocities (yor ini) of 300 km/s. which results in vy, during the Main-Sequence evolution of 180 to 240 km/s. These velocities are within the range observed for O stars(2)."," The rotating models have initial rotational velocities $v_\mathrm{rot~ini}$ ) of 300 km/s, which results in $v_\mathrm{rot}$ during the Main-Sequence evolution of 180 to 240 km/s. These velocities are within the range observed for O stars."776. A table with non-rotating models is provided as Table 5.., A table with non-rotating models is provided as Table \ref{tab:Onorot}.777 As the models only provide a limited mass resolution. a special interpolation routine (described in detail in appendix Appendix A:)) is deployed in order to provide a mass resolution down to | M...," As the models only provide a limited mass resolution, a special interpolation routine (described in detail in appendix \ref{app:evol}) ) is deployed in order to provide a mass resolution down to 1 $M_\odot$."778 For LMC metallicity provide only four models (30. 40. 60 and 120 M... all rotating with 300 km/s). two additional models (15 and 20 M) are taken from the Padova group(?).," For LMC metallicity provide only four models (30, 40, 60 and 120 $M_\odot$, all rotating with 300 km/s), two additional models (15 and 20 $M_\odot$ ) are taken from the Padova group."779. The resulting spectral type mass conversion for LMC metallicity stars is shown in Table 6.., The resulting spectral type mass conversion for LMC metallicity stars is shown in Table \ref{tab:Oz08}.780 For SMC metallicities. only include three (all rotating) models (40. 60 and 120 M..).," For SMC metallicities, only include three (all rotating) models (40, 60 and 120 $M_\odot$ )."781 Again. two models are added here (15 and 20 M..) from in order to derive a spectral type mass conversion (Table 7)).," Again, two models are added here (15 and 20 $M_\odot$ ) from in order to derive a spectral type mass conversion (Table \ref{tab:Oz04}) )."782 The masses shown in the Tables 4. to 7 are all weighted by the duration of the models in each spectral class., The masses shown in the Tables \ref{tab:Orot} to \ref{tab:Oz04} are all weighted by the duration of the models in each spectral class.783 The errors are assigned by using the most- and least-massive model entering the spectral class., The errors are assigned by using the most- and least-massive model entering the spectral class.784 As mentioned before. each spectral class has an assumed error in of 1000 K. The minimal and maximal start and end ages give the range of possible ages for the stars na spectral-class box.," As mentioned before, each spectral class has an assumed error in of 1000 K. The minimal and maximal start and end ages give the range of possible ages for the stars in a spectral-class box."785 The advantage of using this nethod 15 the consistent application of observational constraints for the different evolutionary phases on one set of stellar evolution models., The advantage of using this method is the consistent application of observational constraints for the different evolutionary phases on one set of stellar evolution models.786 This allows one to place more constraints for the stars in a certain spectral class on the range of their initial and present-day masses., This allows one to place more constraints for the stars in a certain spectral class on the range of their initial and present-day masses.787 A number of other mass estimates for spectral types exist in the literature. e.g. and?.. but only the most recent one by MSHOS is used here.," A number of other mass estimates for spectral types exist in the literature, e.g. and, but only the most recent one by MSH05 is used here."788 These models provide spectroscopic stellar masses that are derived from the stellar luminosity. L. and of NLTE stellar atmosphere models through," These models provide spectroscopic stellar masses that are derived from the stellar luminosity, $L$ , and of NLTE stellar atmosphere models through"789 Instituto de sica de Cantabria (CSIC-Universidad de Cantabria). E-39005 Santander. Cantabria. Spain Dpto.,"	 Instituto de sica de Cantabria (CSIC-Universidad de Cantabria), E-39005 Santander, Cantabria, Spain 	Dpto."790 de Fissica Moderna. Universidad de Cantabria. Avda de los Castros s/n. E-39005 Santander. Cantabria. Spain," de Físsica Moderna, Universidad de Cantabria, Avda de los Castros s/n, E-39005 Santander, Cantabria, Spain y nebulae and their central stars is still poorly understood.}"791seen as in Section 5.2.. suggesting the need. for informative priors even for this higher quality data.,"seen as in Section \ref{subsect:sim:current}, suggesting the need for informative priors even for this higher quality data."792 Figure 7 shows the cosmological inference drawn fron the analysis of the mock illSE cluster., Figure \ref{fig:AMI-fgas} shows the cosmological inference drawn from the analysis of the mock iHSE cluster.793 Phe mocdoel-averaged Fahr probability distributions are dominated by the iLISE model contribution this astrophysical model has been successfully. selected. by the evidence., The model-averaged $f_{\rm gas} h$ probability distributions are dominated by the iHSE model contribution – this astrophysical model has been successfully selected by the evidence.794 Taking the median sample as an estimator for the gas fraction we find fu.=(0.061. (512)., Taking the median sample as an estimator for the gas fraction we find $f_{\rm gas}h=(0.061^{+ 0.019}_{- 0.015})$ .795 With no CAIB contamination this estimate changes to fih=(0.054.ULODL011). an Increase. in. precisionMEM of just (rom 28 to 23%)).," With no CMB contamination this estimate changes to $f_{\rm gas}h=(0.054^{+ 0.015}_{- 0.013})$, an increase in precision of just (from 28 to )."796 This is an indication of the small contribution to the error budget that the primordial CM ias al these angular scales., This is an indication of the small contribution to the error budget that the primordial CMB has at these angular scales.797 Indeed. as previously mentioned he shortest ΑΔΗ baseline will be longer than that used in his work. such that the effect of the primordial CM will be reduced: we might therefore expect results Iving inbetween he two situations simulated here.," Indeed, as previously mentioned the shortest AMI baseline will be longer than that used in this work, such that the effect of the primordial CMB will be reduced; we might therefore expect results lying inbetween the two situations simulated here."798 The analysis of the simulated data presented in the previous section was designed to be a simple demonstration of a e&eneral methodology., The analysis of the simulated data presented in the previous section was designed to be a simple demonstration of a general methodology.799 In the current section we discuss the advantages and. disadvantages. of our. approach. and its ability to be extended. beginning with some comparison with other methods currently in use.," In the current section we discuss the advantages and disadvantages of our approach, and its ability to be extended, beginning with some comparison with other methods currently in use."800 The Bayesian. method described here can be seen as a generalisation of the model fitting procedures emploved. by Qier workers., The Bayesian method described here can be seen as a generalisation of the model fitting procedures employed by other workers.801 For example. ο fitted lens models with 1. 2 and 3 parameters to weak shear data for Abell 1689. by the maximum likelihood method: they compute likelihood contours for the parameter uncertainties and. compare the models’ goodness of Lit with the likelihood ratio test.," For example, \citet{GL/KCS02} fitted lens models with 1, 2 and 3 parameters to weak shear data for Abell 1689 by the maximum likelihood method; they compute likelihood contours for the parameter uncertainties and compare the models' goodness of fit with the likelihood ratio test."802 As explained in Section 2.. such a grid-based. computation is not practical with the 6S parameters used. here.," As explained in Section \ref{sect:infer}, such a grid-based computation is not practical with the 6–8 parameters used here."803 Indeed. numerical maximisation of a function of 6S variables is already a demanding problem. especially when there are multiple maxima to be investigated.," Indeed, numerical maximisation of a function of 6–8 variables is already a demanding problem, especially when there are multiple maxima to be investigated."804 Comparing models bv their maximum likelihoods is also rather sensitive to noise features in the data. rather than assessing the relative appropriateness of the models to the task of explaining the data.," Comparing models by their maximum likelihoods is also rather sensitive to noise features in the data, rather than assessing the relative appropriateness of the models to the task of explaining the data."805 The maximum likelihood ratio is formally. equivalent to the evidence ratio when the prior pdf is a delta-function at the best-fit point: this is clearly not an accurate representation of our prior knowledge., The maximum likelihood ratio is formally equivalent to the evidence ratio when the prior pdf is a delta-function at the best-fit point: this is clearly not an accurate representation of our prior knowledge.806 The x fitting of SZ visibility data by 2.. and the joint maximum likelihood analvsis of SZ and X-ray data of ? are two other examples of à small number of parameters being fitted within the context of a single cluster model: one of the aims of this work was to provide a complete framework which combined and extended analyses such as these.," The $\chi^2$ fitting of SZ visibility data by \citet{SZ/Gre++01}, and the joint maximum likelihood analysis of SZ and X-ray data of \citet{J/Ree++02} are two other examples of a small number of parameters being fitted within the context of a single cluster model; one of the aims of this work was to provide a complete framework which combined and extended analyses such as these."807 Other methods proposed for use in the joint analysis of cluster data have. to date. been focused on “parameter[rece reconstruction.," Other methods proposed for use in the joint analysis of cluster data have, to date, been focused on “parameter–free” reconstruction."808 ? suggests an iterative procedure for refining the cluster potential using N-rav.. SZ and weak lensing images. whilst? provide a direct inversion method to take these images and produce a threc-dimensional cluster moclel. making use of some attractive features of working in the Fourier domain.," \citet{J/Reb00} suggests an iterative procedure for refining the cluster potential using X-ray, SZ and weak lensing images, whilst \citet{J/Zar++01} provide a direct inversion method to take these images and produce a three-dimensional cluster model, making use of some attractive features of working in the Fourier domain."809 7 suggest using SZ and weak lensing cata to constrain successive orders of perturbation from spherical symmetry. again working from ready-made maps.," \citet{J/Dor++01} suggest using SZ and weak lensing data to constrain successive orders of perturbation from spherical symmetry, again working from ready-made maps."810 ALL these methods have been shown to work well with noiseless input., All these methods have been shown to work well with noiseless input.811 Llowever. those working with the observations have opted to increase the complexity of their cluster models in a more eraclual wav. for example fitting an N-rav map with a simple model and using this model to predict. the observed SZ ellect (e.g.2)..," However, those working with the observations have opted to increase the complexity of their cluster models in a more gradual way, for example fitting an X-ray map with a simple model and using this model to predict the observed SZ effect \citep[\eg][]{J/Jon++01}."812 The method described here is this common sense reduced to caleulation the use of simple functions [or the potential. gas density ancl temperature allows the model complexity to be tuned to the data quality via the evidence.," The method described here is this common sense reduced to calculation – the use of simple functions for the potential, gas density and temperature allows the model complexity to be tuned to the data quality via the evidence."813 “Phe 7parameter-free methods referred. to above actually have many parameters. usually the values of pixels in à grid: clirect methods will produce one set of parameters. but these may not be the most probable given the data. or even the most appropriate given the parameter degeneracies.," The “parameter-free” methods referred to above actually have many parameters, usually the values of pixels in a grid: direct methods will produce one set of parameters, but these may not be the most probable given the data, or even the most appropriate given the parameter degeneracies."814 When trving to measure quantities such as the gas fraction. all the cluster configurations allowed by the data should be accounted. for. in order to calculate an accurate confidence interval: only by Lully exploring a models parameter space this can be achieved.," When trying to measure quantities such as the gas fraction, all the cluster configurations allowed by the data should be accounted for, in order to calculate an accurate confidence interval; only by fully exploring a model's parameter space this can be achieved."815 Indeed. consideration of a range of models is then desirable to eain the next level of accuracy. one that is mocel-independent in the sense of the discussion in Section 2..," Indeed, consideration of a range of models is then desirable to gain the next level of accuracy, one that is model-independent in the sense of the discussion in Section \ref{sect:infer}."816 Sampling from a models parameter space produces the set. of. eluster. configurations permitted by the data. which is arguably more useful than the unique solutions generated by direct methods.," Sampling from a model's parameter space produces the set of cluster configurations permitted by the data, which is arguably more useful than the unique solutions generated by direct methods."817 One might wonder how the ALCALC technique endorsed in this work would perform if used in the many-parameter modelling of the type mentioned above., One might wonder how the MCMC technique endorsed in this work would perform if used in the many-parameter modelling of the type mentioned above.818 Indeed. in this context. the number of parameters. included: here is rather small.," Indeed, in this context the number of parameters included here is rather small."819" The evidence itself is a guide in the development of. these methods. providing a handle. on the information content of the data: if the evidence does not favour a triaxial cllipsoid over a spherical model. then it nueht reasonably be assumed not to favour a 7""parameter-free representation."," The evidence itself is a guide in the development of these methods, providing a handle on the information content of the data: if the evidence does not favour a triaxial ellipsoid over a spherical model, then it might reasonably be assumed not to favour a “parameter-free” representation."820 The ability of a sampler to cope with increasing numbers of parameters is somewhat sensitive to the shape of the posterior cistribution under investigation: we find that including extra nuisance parameters (the parameters of point sources contaminating the SZ data for example) does not alfect the accuracy of the posterior exploration (Lancaster 2003 in preparation). and neither does increasing. the number of sub-clumps when modelling gravitational lenses with multiple mass concentrations (7)..," The ability of a sampler to cope with increasing numbers of parameters is somewhat sensitive to the shape of the posterior distribution under investigation; we find that including extra nuisance parameters (the parameters of point sources contaminating the SZ data for example) does not affect the accuracy of the posterior exploration (Lancaster 2003 in preparation), and neither does increasing the number of sub-clumps when modelling gravitational lenses with multiple mass concentrations \citep{GL/Kne++03}. ."821 Moving to three-dimensional cluster modelling introduces a number of strong degeneracies in the parameter space (2). which may well require a tailor- ALCAIC sampler instead of thegeneral purpose engine, Moving to three-dimensional cluster modelling introduces a number of strong degeneracies in the parameter space \citep{J/F+P02} which may well require a tailor-made MCMC sampler instead of thegeneral purpose engine822 (SDSS.Yorketal.2000).," \citep[SDSS,][]{York2000}."823.. Wavetal.(2009.hereafterPaperII). Petrosiau (e.g.Collister&Lahay2008).," \citet[][hereafter Paper II]{Way09} \cite{Petrosian1976} \citep[e.g.][]{CL2004,Ball2004,Wadadekar05,Kurtz2007,WG08}."824. Singaletal.(2011). Yipetal.(2011) of disk ealaxies., \cite{Singal2011} \cite{Yip2011} of disk galaxies.825 Ou the other hand. attempts to morphologically classify lavee number of galaxies iu the universe has eained inaccuracy over the past 15 vears as better/lareer {παλιο samples from eve classification has increased.," On the other hand, attempts to morphologically classify large number of galaxies in the universe has gained in accuracy over the past 15 years as better/larger training samples from eye classification has increased."826 For example. Lahawetal.(1995) was one of the first to use an Artificial Neural Network trained on 830 ealaxies classified by the eves of six different professional astronomers.," For example, \cite{Lahav1995} was one of the first to use an Artificial Neural Network trained on 830 galaxies classified by the eyes of six different professional astronomers."827 Iu more recent vears Balletal.(2001) has attempted to classify galaxies by morphological type using a Neural Network approach based on a sample of 1399 ealaxies (frou the catalog of Nakamuractal. (2, In more recent years \cite{Ball2004} has attempted to classify galaxies by morphological type using a Neural Network approach based on a sample of 1399 galaxies (from the catalog of \cite{Nakamura2003}) ).828011))). Chengetal.(2011). has used a sample of 981 uon-star forming SDSS early-type galaxies to distinguish between E. S0 aud Sa galaxies.," \cite{Cheng2011} has used a sample of 984 non-star forming SDSS early-type galaxies to distinguish between E, S0 and Sa galaxies."829 In the past year two new attempts at morphological classification using Machine Learning techuiques ou a Galaxy Zoo (Lintottet2011) traimineg sample have been published (Banerjictal.20100:Tnertas-Companyet 2011).," In the past year two new attempts at morphological classification using Machine Learning techniques on a Galaxy Zoo \citep{Lintott2008,Lintott2011} training sample have been published \citep{Banerji2010,HC2011}."830. The Banerjiotal.(2010) results were impressive in that thev claim to obtain classification to better than for three different morphological classes (spiral. elliptical and poiut-sources/artifacta).," The \cite{Banerji2010} results were impressive in that they claim to obtain classification to better than for three different morphological classes (spiral, elliptical and point-sources/artifacts)."831 These works are iu contrast to previous work like that of Bernardietal.(2003) who used a classification scheme based on SDSS spectra.," These works are in contrast to previous work like that of \cite{Bernardi2003}832 who used a classification scheme based on SDSS spectra."833 However. this classification certainly iudssed some carly-type salaxies from their desired sauple due to the preseuce of star formation.," However, this classification certainly missed some early-type galaxies from their desired sample due to the presence of star formation."834 Tn this paper we will coutinue our use of Cassia Process Reeression to calculate Photo-Zs. using a varietyof iuputs.," In this paper we will continue our use of Gaussian Process Regression to calculate Photo-Zs, using a varietyof inputs."835 This method has been discussed extensively in two previous papers (Way&Srivastava2006:Way 20049)..," This method has been discussed extensively in two previous papers \citep{Way06,Way09}. ."836 We utilizetheSDSS Main Galaxy Sample (AIGS. and the Lunuüuous Red Calaxy Sample (LRG.Eiseusteiuetal.2001). from the SDSS Data Release Seven (DR7.Abazajianetal. 2009)...," We utilizetheSDSS Main Galaxy Sample \citep[MGS,][]{Strauss02} and the Luminous Red Galaxy Sample \citep[LRG,][]{Eisenstein01} from the SDSS Data Release Seven \citep[DR7,][]{SDSS07}. ."837 We, We838We can eliminate the pressure variableIT from equations (19) and (21) by cross differentiation. reducing the svstem to one second order equation. equation (20). and one fourth order equation. given by There are two distinct classes of solutions to the svstem (20).(22).,"We can eliminate the pressure variable$\Pi$ from equations (19) and (21) by cross differentiation, reducing the system to one second order equation, equation (20), and one fourth order equation, given by There are two distinct classes of solutions to the system (20),(22)."839 Since these equations are linear. these solutions can be superimposed with relative amplitudes determined by the boundary conditions for each.," Since these equations are linear, these solutions can be superimposed with relative amplitudes determined by the boundary conditions for each."840 One class is found by setting the meridional How streaamfunction io—0 evervwhere., One class is found by setting the meridional flow streamfunction $\psi=0$ everywhere.841 It. vields pure differential rotation independent of z forced at the outer boundary of the exlinder., It yields pure differential rotation independent of $z$ forced at the outer boundary of the cylinder.842 To satisfy the condition that there be no net torque at the outer boundary. to allow steady solutions in the interior of the evelinder. there must be no viscous stress al this boundary.," To satisfy the condition that there be no net torque at the outer boundary, to allow steady solutions in the interior of the cyclinder, there must be no viscous stress at this boundary."843 This corresponds {ο solutions with constant angular velocity. or linear rotational velocity à that decreases linearly with s toward the axis of the evlincler.," This corresponds to solutions with constant angular velocity, or linear rotational velocity $u$ that decreases linearly with $s$ toward the axis of the cylinder."844 The other class of solutions can be found by separation of variables. since coellicients in this svslem are functions of s only.," The other class of solutions can be found by separation of variables, since coefficients in this system are functions of $s$ only."845 We place the lower ancl upper boundaries of our evlinder αἱ zo—0.H respectively.," We place the lower and upper boundaries of our cylinder at $z=0,H$ respectively."846 Then if we allow no flow through the lower or upper boundary. these boundaries must coincide with a streamline. and there should be no viscous stress (here either.," Then if we allow no flow through the lower or upper boundary, these boundaries must coincide with a streamline, and there should be no viscous stress there either."847 These conditions are satisfied if we take With this choice of representation of the solutions. we can find separate solutions for each n for i and 4.," These conditions are satisfied if we take With this choice of representation of the solutions, we can find separate solutions for each $n$ for $\psi$ and $u$."848" We can then represent the forcing at the boundary in thesame war. so the amplitude of the solutions for each η is determined separately by the amplitude of the forcing [ον the same η,"," We can then represent the forcing at the boundary in thesame way, so the amplitude of the solutions for each $n$ is determined separately by the amplitude of the forcing for the same $n$."849 There could be other solutions that use other representations. bul we have not looked for them.," There could be other solutions that use other representations, but we have not looked for them."850" Then if we substitute expressions (23) into equations (20).(22). and define o,=nz/H. then equations (20) and (22) become. for each n. and in which"," Then if we substitute expressions (23) into equations (20),(22), and define $\sigma_n=n\pi/H$ , then equations (20) and (22) become, for each $n$ , and in which"851erant AST 09-08114 to JIIU 1.1.) and by NASA erant NNNOO9AIIISA (The Kepler Open Cluster Study).,grant AST 09-08114 to JHU (I.P.) and by NASA grant NNX09AH18A (The Kepler Open Cluster Study).852 Astrophysics at Keele University is supported by an STFC Rolling Grant., Astrophysics at Keele University is supported by an STFC Rolling Grant.853 This study is based on observations obtained with Megal?rime/MegaCzun. a joint project of Canada-France-Hawali Telescope (CFUT) and CEA/DAPNIA. at the CFIUT which is operated by the NRC of Canada. the Institute National des Sciences de [Universe of the CNRS. and the University of Hawai.," This study is based on observations obtained with MegaPrime/MegaCam, a joint project of Canada-France-Hawaii Telescope (CFHT) and CEA/DAPNIA, at the CFHT which is operated by the NRC of Canada, the Institute National des Sciences de l'Universe of the CNRS, and the University of Hawaii."854perpendicular to the orbital plane.,perpendicular to the orbital plane.855" Quantitatively, at the end of the simulation the mass of gas at large distances (|z|>10 kpc) from the orbital plane is 3x10*?M in the simulation including the AGN wind, about 20 times larger than in the simulation with only radiation pressure feedback (see Table 1 and Fig."," Quantitatively, at the end of the simulation the mass of gas at large distances $|z| > 10$ kpc) from the orbital plane is $\simeq 3856\times 10^9 \msun$ in the simulation including the AGN wind, about 20 times larger than in the simulation with only radiation pressure feedback (see Table \ref{tab:simparm} and Fig."857 6 below)., \ref{fig:woblowout} below).858" A corollary of this efficient removal of gas by the AGN wind is that the total stellar mass formed during the simulation is ~20% smaller in the case with the AGN wind; most of this suppression in star formation happens at late times, after the final coalescence of the two galaxies (see Fig. 1))."," A corollary of this efficient removal of gas by the AGN wind is that the total stellar mass formed during the simulation is $\sim 20 \%$ smaller in the case with the AGN wind; most of this suppression in star formation happens at late times, after the final coalescence of the two galaxies (see Fig. \ref{fig:mdotfid}) )."859" 'The evacuation of the central part of the galaxy by the AGN wind would not be surprising if the majority of the material ejected was explicitly added to the wind by our deposition of momentum (i.e., if the unbound mass was primarily material that was explicitly ’kicked’)."," The evacuation of the central part of the galaxy by the AGN wind would not be surprising if the majority of the material ejected was explicitly added to the wind by our deposition of momentum (i.e., if the unbound mass was primarily material that was explicitly 'kicked')."860 We find that this is not the case., We find that this is not the case.861" At the end of the fiducial simulation (t=2.85 Gyr), of the 33,028 gas particles at large transverse positionkpc), only 8,904 of the particles have been explicitly kicked."," At the end of the fiducial simulation $t = 2.85$ Gyr), of the 33,028 gas particles at large transverse position, only 8,904 of the particles have been explicitly kicked."862 The other by mass) have been ejected because of hydrodynamic interactions with wind material., The other by mass) have been ejected because of hydrodynamic interactions with wind material.863" In addition, at any time only a small fraction ~5% of the gas inside Race (the accretion/feedback region) has been explicitly added to the wind."," In addition, at any time only a small fraction $\sim 5 \%$ of the gas inside $R_{acc}$ (the accretion/feedback region) has been explicitly added to the wind."864 These results demonstrate that the majority, These results demonstrate that the majority865The wmuber of spectroscopic observations of supernovae increased quickly in this decade. partly due to the advances of CCD-techuique and the erowine nuniber of supernova searchlius observational projects (see Filippeuko. 1997. for a detailed review).,"The number of spectroscopic observations of supernovae increased quickly in this decade, partly due to the advances of CCD-technique and the growing number of supernova searching observational projects (see Filippenko, \cite{filip} for a detailed review)."866 In the first half of 1998. three SNe with brightuess of about Vz12 mag (199558. 1998aq aud 1998bu) was discovered together with many niore fainter ones.," In the first half of 1998, three SNe with brightness of about $V{\approx}12$ mag (1998S, 1998aq and 1998bu) was discovered together with many more fainter ones."867 Iu this paper we report iiediun-resolutiou spectroscopic observations of SN 1998aq made shortly before aud after maxinuun lieht., In this paper we report medium-resolution spectroscopic observations of SN 1998aq made shortly before and after maximum light.868 SN L998aq was discovered. by. AL. Armstrong αντ et al. 1998)).," SN 1998aq was discovered by M. Armstrong (Hurst et al., \cite{hurst}) )."869 It has been classified as SN Ia by Avani Yamaoka (1998)) who reported promineu Si II AG355 aud other S IL. Fe IT and Ale II absorption lines which mace SN 1998aq siuuilar to the “prototype” SN Ta SN 1991D. The expansion velocity was determined as about 11.000 kun/s. They also poiuted out the absence of Na D absorption due to probably simall iuterstellar reddening.," It has been classified as SN Ia by Ayani Yamaoka \cite{ayani}) ) who reported prominent Si II $\lambda$ 6355 and other S II, Fe II and Mg II absorption lines which made SN 1998aq similar to the “prototype” SN Ia SN 1994D. The expansion velocity was determined as about 11,000 km/s. They also pointed out the absence of Na D absorption due to probably small interstellar reddening."870 Shortly later. Berlind Calkins (see Carnavich et al. 1998))," Shortly later, Berlind Calkins (see Garnavich et al., \cite{garnav}) )"871 reported the similarity to SN 1990N. based on a spectrum obtained at 1 week before maxima., reported the similarity to SN 1990N based on a spectrum obtained at 1 week before maximum.872 Another interesting property of SN 1998aq is that its host ealaxy. NGC 3982 (PCC 237520. UCC 69185. TRAS 1153815521). has a Sevfert 2 type nucleus.," Another interesting property of SN 1998aq is that its host galaxy, NGC 3982 (PGC 37520, UGC 6918, IRAS 11538+5524), has a Seyfert 2 type nucleus."873 This ealaxy was a subject of a recent. ACN-survey by Πο et al. (1997))., This galaxy was a subject of a recent AGN-survey by Ho et al. \cite{ho}) ).874 To date. there is an indication that SNe in the host galaxies of ACNs show higher couceutration toward the galaxw corewith respect to SNe in normal galaxies (Petrosian Turatto. 1990)).," To date, there is an indication that SNe in the host galaxies of AGNs show higher concentration toward the galaxy corewith respect to SNe in normal galaxies (Petrosian Turatto, \cite{petros}) )."875" This may give evidence ou increased star formation rate in the proximity of ACN, but the umber of actually observed SNe in such svsteuis is not large. so more data would significantly improve the statistics."," This may give evidence on increased star formation rate in the proximity of AGN, but the number of actually observed SNe in such systems is not large, so more data would significantly improve the statistics."876 We made medi and high-resolution spectroscopic observations of SN 1998aq between April 22th and Ma 27th. 1998 at David Dunlap Observatory. Canada with he 7 Casscerain telescope.," We made medium- and high-resolution spectroscopic observations of SN 1998aq between April 22th and May 27th, 1998 at David Dunlap Observatory, Canada with the 74"" Cassegrain telescope."877 The eratinges used were he 150 lines/uuu Gn 2ud order with an order-separation filter inserted) and the P800 Πποςπμ giviug 1.3 per κο] and 0.2 per pixel resolution. respectively.," The gratings used were the 150 lines/mm (in 2nd order with an order-separation filter inserted) and the 1800 lines/mm giving 1.3 per pixel and 0.2 per pixel resolution, respectively."878 The uediuu-esolutiou spectra are presented in Fie., The medium-resolution spectra are presented in Fig.879 d (loft uel) where an arbitrary vertical shift has been added o each spectrum for better visbilitv., \ref{fig_1} (left panel) where an arbitrary vertical shift has been added to each spectrum for better visibility.880 The decrease of he signal-to-noise ratio toward the later spectra was due to the fünutenimg aud the increasing airmass of the object in Max., The decrease of the signal-to-noise ratio toward the later spectra was due to the faintening and the increasing airmass of the object in May.881 The data were reduced by standard[RAF routines., The data were reduced by standard routines.882 FeAr spectral lamp exposures were used for wavelength calibration., FeAr spectral lamp exposures were used for wavelength calibration.883 Particular attention was paved to remove the backeround Πο coutanünatiou due to the host galaxy (discussed below) aud the night sk., Particular attention was payed to remove the background light contamination due to the host galaxy (discussed below) and the night sky.884 Au unfiltered CCD-inaee showing SN 1998aq in NGC 3982 taken from downtown of Szeged with a 11 Schinidt- telescope and ST-6 camera is prescuted in the vielt panel of Fig. 1..," An unfiltered CCD-image showing SN 1998aq in NGC 3982 taken from downtown of Szeged with a 11"" Schmidt-Cassegrain telescope and ST-6 camera is presented in the right panel of Fig. \ref{fig_1}."885 Iu order to determine the phase of our spectra relative to the light curve of the SN. we collected all. available visual observatious of SN 1998aq made bv amateur astronomers. using the public database of the (," In order to determine the phase of our spectra relative to the light curve of the SN, we collected all available visual observations of SN 1998aq made by amateur astronomers, using the public database of the ."886VSNET) This light curve is, This light curve is887The formation of the Milky Way (MW) galaxy is a mystery unsolved yet.,The formation of the Milky Way (MW) galaxy is a mystery unsolved yet.888 Different models are trying to explain what were the initial conditions that lead to the actual structure of the MW., Different models are trying to explain what were the initial conditions that lead to the actual structure of the MW.889 It is commonly accepted that the structure of the Milky Way. and other comparable disc galaxies. can be divided into three main components. the bulge. the galactic spheroid and the disc.," It is commonly accepted that the structure of the Milky Way, and other comparable disc galaxies, can be divided into three main components, the bulge, the galactic spheroid and the disc."890 The central bulge has a mass of zz107 M... and a characteristic radius of about | kpe., The central bulge has a mass of $\approx 10^{10}$ $_{\odot}$ and a characteristic radius of about $1$ kpc.891 The galactic spheroid. which is also called the stellar halo. has à mass of z3.7cL210 M. (Bell2008) and its mass is mostly contined within the the Solar Radius.," The galactic spheroid, which is also called the stellar halo, has a mass of $\approx 3.7 \pm 1.2 \times89210^{8}$ $_{\odot}$ \citep{bell08} and its mass is mostly confined within the the Solar Radius."893" The stellar halo contains globular clusters. dwarf satellites and their tidal streams. that add up to a mass of about LO?"" M..."," The stellar halo contains globular clusters, dwarf satellites and their tidal streams, that add up to a mass of about $10^{6-7}$ $_{\odot}$."894" For the MW. the dise ean be subdivided into at least two parts: the thin dise with a mass of about Aq,=5«107 M.. that has exponential radial and vertical scale lengths of approximately hy=2.3+0.6 kpe (Hammeretal.2007) and hf,z&300 pe (Juricetal.2008). respectively."," For the MW, the disc can be subdivided into at least two parts: the thin disc with a mass of about $M_{\rm disc} = 5 \times89510^{10}$ $_{\odot}$ that has exponential radial and vertical scale lengths of approximately $h_{\rm R} = 2.3 \pm 0.6$ kpc \citep{hammer07}896 and $h_{\rm z} \approx 300$ pc \citep{jur08} respectively."897" The other part is the thick disk which has scale lengths of μια=4.1+0.4 kpe and fina,=0.75+0.07 kpe (deJongetal.2010)."," The other part is the thick disk which has scale lengths of $h_{\rm thd,R} = 4.1 \pm 0.4$ kpc and $h_{\rm thd,z}898= 0.75 \pm 0.07$ kpc \citep{jong10}."899 Near the Sun. the thick dise comprises about 6 per cent of the thin disc mass. so that the thick dise mass amounts to δημz02.0:55Mais.," Near the Sun, the thick disc comprises about $6$ per cent of the thin disc mass, so that the thick disc mass amounts to $M_{\rm thd} \approx 0.2-0.3 \times M_{\rm disc}$."900" The thick disc is made up mostly of low-metalicity (e/14]:0.4) stars that have a velocity dispersion perpendicular to the dise plane of 0,5.7240 . compared to the significantly smaller c of the thin disc. which varies from about 2.05 |! for the young stars (Fuchsetal. 2001).."," The thick disc is made up mostly of low-metalicity $([Fe/H] \leq -0.4)$ stars that have a velocity dispersion perpendicular to the disc plane of $\sigma_{\rm z,obs} \approx 40$ $^{-1}$, compared to the significantly smaller $\sigma_{\rm z}$ of the thin disc, which varies from about $2-5$ $^{-1}$ for the young stars \citep{Fuchs2001}. ."901 These authors measured the velocity dispersion in the solar neighbourhood as a function of age of the stars., These authors measured the velocity dispersion in the solar neighbourhood as a function of age of the stars.902 The oldest stars in their sample (CNS4 have about 25 peMMyr+ for 10 Gyr old stars., The oldest stars in their sample (CNS4) have about $25$ $^{-1}$ for $10$ Gyr old stars.903" But then this value might be ""contaminated! by thick dise stars.", But then this value might be 'contaminated' by thick disc stars.904 Several mechanisms have been proposed to explain the formation of the thick dise in galaxies., Several mechanisms have been proposed to explain the formation of the thick disc in galaxies.905 One of these mechanisms was proposed by Abadietal.(2003)... who suggest that the formation of the thick dise is the direct accretion of stars from disrupted satellites.," One of these mechanisms was proposed by \citet{aba03}, who suggest that the formation of the thick disc is the direct accretion of stars from disrupted satellites."906 The process of accretion occurs approximately at coplanar orbits., The process of accretion occurs approximately at coplanar orbits.907 Another explanation was suggested by RoSkaretal.(2008) and Schónrich&Binney(2009).. who consider the process of radial migration of the stars.," Another explanation was suggested by \citet{ros08} and \citet{scho09}, who consider the process of radial migration of the stars."908 In this mechanism. the stars which end up in the thick dise are trapped onto a resonant co- with spiral arms and may migrate inwards and outwards along the spiral waves.," In this mechanism, the stars which end up in the thick disc are trapped onto a resonant co-rotation with spiral arms and may migrate inwards and outwards along the spiral waves."909 This process conserves angular momentum and does not lead to significant heating of the disc., This process conserves angular momentum and does not lead to significant heating of the disc.910 Another »xossible scenario is proposed by citetquin93.. Kazantzidisetal.(2008). and Villalobos&Helmi(2008) and consists of the thickening of a pre-existing thin disc hrough minor mergers.," Another possible scenario is proposed by \\citet{quin93}, , \citet{kaz08} and \citet{villa08} and consists of the thickening of a pre-existing thin disc through minor mergers."911 The thick dise is formed by the dynamical yeating that is induced by satellites merging with a primordial. rotationally supported thin disc.," The thick disc is formed by the dynamical heating that is induced by satellites merging with a primordial, rotationally supported thin disc."912 Finally.Brooketal.(2005). and Bornaudetal.(2007) suggest that the formation of the thick dise isriggered in situ.," Finally,\citet{brok05} and \citet{bour07} suggest that the formation of the thick disc istriggered in situ."913 The process of star formation occurs during/aftergus rich mergers., The process of star formation occurs during/aftergas rich mergers.914 Each model explains different aspects and has its own, Each model explains different aspects and has its own915 Fic. 4.—Sideandtop viewsofa snapshotof mag,"previously accumulated by the system (Figure \ref{fig1}) ).,"916netic field-lines has been rescaledforan improved visualization., which would otherwise bring the system to a configuration similar to the initial condition at $t=0$.917Top: Fieldlines of(z—5) field-linesof the velocityorthogonal componenttopplateof(z the squaredboundarycurrentforcingj?. isosu," This dissipative event is in fact due to magnetic reconnection, that during its evolution produces a component of the magnetic field along $x$, the cross-shear direction, forming magnetic islands [see Figure \ref{fig3} at times $t \sim 79$ and $82\, \tau_A$ ]."918"rfaceinthe j?—10), the10? the isosurface withj?=8 x The10°,well belowatthe maximumvalue 2.8ofxthe isos"," In fact around $t \sim 90\, \tau_A$, at the end of the big dissipative event, the topology of the orthogonal component of the magnetic field is characterized by magnetic islands."919urfaces corresponding to higher values, Naturally the Lorentz force does not vanish now and the vorticity is not constant along the streamlines.920of j? are nestedinside small., As typical of magnetic reconnection vorticity forms misaligned quadrupolar structures around current sheets [see \cite{rved08}] ].921 forcingvelocity at the boundary would recreate overtime ," Although the forcing velocity at the boundary is always a shear [eqs. \ref{eq:f0})\ref{eq:f1}) )],"922a sheared magnetic fieldinthe syst," as they do during the linear stage for $t < 79\, \tau_A$."923em that shouldthen leadto anotherbigdissipative event andsoon., When they vanish magnetic energy can be stored without getting dissipated (see \\ref{par3}) ).924" The dynamicsof this system are infact commonly approximated asa sequence of equilibria, each destabilized by magnetic reconnection."," But now nonlinearity can along the cross-shear $x$ ) direction part of the energy associated with the shear-aligned $y$ -oriented) field along which the forcing injects energy, and continuously cascades to lower scales as described in the following sections."925 This approximation is , The three-dimensional structures are shown in Figure \ref{fig4}.926attained by neglecting the velocity and kinetic pressurein theMHD equations whose solu," Although the magnetic energy dominates over the kinetic energy, the ratio of the rms of the orthogonal magnetic field over the axial dominant field $B_0$ is quite small."927"tion is then bound tobea static force-free equilibrium. simulations the system is magnetically dominated, and in particular magnetic"," For $c_A = 200$ it is $\sim 3\%$, so that the average inclination of the magnetic field-lines with respect to the axial direction is just $\sim 2^{\circ}$, it is only for lower value of $c_A$ that this ratio increases and the angle increases accordingly \citep{rved08}."928" energyis bigger than kinetic energy, asshown in Figure [Il] "," The field-lines of the total magnetic field at time $550\, \tau_A$ are shown in Figure \ref{fig4} )."929whereonthe averageEy 61 Ej. Butthe self-consistent evolution ofthekin," The computational box has been rescaled for an improved viewing, and to attain the original aspect ratio, the box should be stretched 10 times along the axial direction."930"etic pressure and velocity, alt"," The magnetic topology for the total field is quite simple, asthe lines appear slightly bent."931hough small row) and current sheets (bottomrow) attim," Figure \ref{fig4} ) also shows a view from the side and the top of the 3D current sheets at time $550\, \tau_A$."932"eT~ 55074.The box Intotalthe magneticsidefield have (orthogonalsuperimposedplus axial),yello"," The current sheets, elongated along the axial direction, look space filling when watched from the side of the computational box, but the view from the top shows that the filling factor is actually small, as they are almost 2D structures."933w andinthe streamlinesmidplane we, So far we have analyzed the topology of the field-lines only in the mid-plane $z=5$.934 imposeviewawe vanishing velocity. in 'Twoisosurfaces bottom representedin ," In Figure \ref{fig5} we show the current density and magnetic field lines of $\mathbf{b}_{_\perp}$ in the mid-plane $z=5$ ) and in two other $x$ $y$ planes close to the boundaries $z=1$ and $9$, the axial length $L=10$ )."935"transparent yellow,Bottom:whilered displays theof current at this time Ίρις partially=3.6x107."," The behavior is similar at different heights although in the plane $z=9$, closer to the forced boundary $z=10$, the topology of the field appears to be affected to some extent by the sheared velocity forcing \ref{eq:f0}) ) directed along the $y$ direction."936 Asistypical of current she, The field-lines in fact show a small alignment directed along $y$ close to the boundary $z=10$ .937"ets, corresponding tolower values. Thecurrent sheetsfilling factoris comparedwith the dominantaxial magneticfield", The influence of the boundary forcing over the magnetic field can be expressed through the correlation between the magnetic field $\mathbf{b_{_\perp}}$ in the plane $z$ and the boundary forcing velocity $\mathbf{u^L}$ \ref{eq:f0}) )]: In Figure \ref{fig6} we plot the correlation as a function of the axial coordinate $z$ at selected times.938" Bo, does not bindthe system to force-free equilibria allows the possible development of alternative dynamics."," In the linear stage, until time slightly bigger than $t = 76\, \tau_A$ whereafter the system transitions to the nonlinear stage (Figures \ref{fig1}, \ref{fig2}, and \ref{fig3}) ), the magnetic field is a mapping of the boundary velocity therefore as expected the correlation is equal to 1."939 Inthe following sectionswe will analyze further aspectsof the dynamicsand the spectral properties ofthe system. Butfirst ," In fact for the simulation presented in this section (run A), for which we have imposed the shear velocity profile \ref{eq:f0}) ) at the top plate $z=10$ and a vanishing velocity at the bottom plate $z=0$, the magnetic field in the linear stage is given by eq. \ref{eq:lin1s}) )["940we illustrate the topologyof the magnetic and velocityf,or \ref{eq:diff1}) ) with $\mathbf{u^0}=0$ including diffusion] therefore the correlation is 1 as $\mathbf{b_{\perp}}$ is proportional to the boundary velocity $\mathbf{u^L}$.941"ields, to understand whya sheared magnetic fieldisnotrecreated. MagneticField Topology and Originof Turbulence the Asshownin Figure3] attime t~79TA re"," Next as the system transitions to the nonlinear stage releasing most of the accumulated magnetic energy the correlation between the magnetic field and the boundary forcing velocity decreases swiftly, at a faster pace the farther from the forced boundary $z=10$ , as shown by the curves at times $82.19 \le t / \tau_A \le 88.28$ ."942"connection startsto develop, enhancing the ohmic dissipation, that reachesa peak aroundt ~ 8274."," The correlation during the fully nonlinear stage is shown with color lines at $10$ selected times separatedby $\Delta\, t = 40\, \tau_A$ in the interval $200\, \tau_A \le t \le 600\, \tau_A$ ."943 Thisbig dissipative event burns alarge fraction, The correlation vanishes near the bottom boundary and then grows almost linearly with $z$ up to $\sim 0.6$ at the top boundary.944ofthe magnetic energy, As expected the correlation is bigger near the945various epochs.,various epochs.946 To identify groups we link pairs of infalling halos whose angular momentum orientations are separated by a<10 and with relative distances ef«40 kpe at the time of accretion.," To identify groups we link pairs of infalling halos whose angular momentum orientations are separated by $\alpha <94710^\circ$ and with relative distances $d < 40$ kpc at the time of accretion."948 We found that this combination of à and d values results in a robust set of groups. maximizing their extent while minimizing the number of spurious links.," We found that this combination of $\alpha$ and $d$ values results in a robust set of groups, maximizing their extent while minimizing the number of spurious links."949 We follow the orbits ofthe groups identified from redshift until present time.," We follow the orbits of the groups identified from redshift $z \sim9504.2$ until present time."951 Fig., Fig.952 3 shows the trajectories of some of the richest groups of subhalos. which were accreted 2.43. 1.65 and 0.84 Gyrs ago respectively.," \ref{orbital_plot} shows the trajectories of some of the richest groups of subhalos, which were accreted 2.43, 1.65 and 0.84 Gyrs ago respectively."953 Each dot represents the position of a subhalo colour coded from high-redshift (dark) to the present (light-grey)., Each dot represents the position of a subhalo colour coded from high-redshift (dark) to the present (light-grey).954 The erosses correspond to the present«lay positions while those at the time of accretion are shown as open circles., The crosses correspond to the present-day positions while those at the time of accretion are shown as open circles.955 Fig., Fig.956 3. clearly shows that the groups of subhalos follow nearly coherent orbits as early as 2~4.2. long before the time of accretion.," \ref{orbital_plot} clearly shows that the groups of subhalos follow nearly coherent orbits as early as $z \sim9574.2$, long before the time of accretion."958 The characteristic size of the groups can be measured by computing the number function of groups. i.e. how many groups have a given number of subhalos.," The characteristic size of the groups can be measured by computing the number function of groups, i.e. how many groups have a given number of subhalos."959 Fig., Fig.960 + shows the number function of groups accreted in the four most recent snapshots: present time. 0.84. 1.65. and 2.43 Gyrs ago.," \ref{grp_richness} shows the number function of groups accreted in the four most recent snapshots: present time, 0.84, 1.65, and 2.43 Gyrs ago."961 As can be seen from this Figure. the shape is quite similar at all times. and most of the groups have a small number of members.," As can be seen from this Figure, the shape is quite similar at all times, and most of the groups have a small number of members."962 Fig., Fig.963 5 shows the differential mass function of the groups in Fig., \ref{massfn_allgrp} shows the differential mass function of the groups in Fig.964" 4. down to our resolution limit (the dashed line. which corresponds to 5.89«10""AL. y."," \ref{grp_richness} down to our resolution limit (the dashed line, which corresponds to $\sim 5.89 \times 10^{6}\msun$ )."965 Once again we find very similar power-law shapes for the mass functions at different epochs., Once again we find very similar power-law shapes for the mass functions at different epochs.966 This power-law shape is reminiscent of the differential mass function of subhalos in cluster and galaxy-size dark matter halos., This power-law shape is reminiscent of the differential mass function of subhalos in cluster and galaxy-size dark matter halos.967" The slope of the fitted dNfdlogMxM"" relation is n~0.5+0.2.", The slope of the fitted $dN/d\log M \propto M^{n}$ relation is $n \sim -0.5 \pm 0.2$.968 Note hat this is somewhat shallower than that found for subhalos. where noceOS+0.1 (Stoehretal.2003:DeLucia2004:Gaoal. 2004b).," Note that this is somewhat shallower than that found for subhalos, where $n\sim -0.8 \pm 0.1$ \citep{stoehr03,delucia04,gao04b}."969". This could well be due to insufficient mass resolution: he fact that we are not resolving subhalos below 2.910""M... implies that many subhalos are accreted in isolation. instead of in xuirs or in groups."," This could well be due to insufficient mass resolution: the fact that we are not resolving subhalos below $2.9970\times 10^6 \msun$, implies that many subhalos are accreted in isolation, instead of in pairs or in groups."971 This effect is much stronger at the low mass end of the group mass spectrum., This effect is much stronger at the low mass end of the group mass spectrum.972" For example. a group with total mass ~JO""M. ean consist of ten subhalos of ~LO“M. or two of"," For example, a group with total mass $\sim 10^9 \msun$ can consist of ten subhalos of $\sim 10^8 \msun$ or two of"973For the fidueial parameters. {μεcfee)2107 em.,"For the fiducial parameters, $R_{\rm in}(t>t_{\rm dec})\sim 2\times 10^{17}$ cm."974" The strengthe of the magnetico field at the inner radius By, can be estimated as D,2«10.U C. which is almost the same as that of ISAL"," The strength of the magnetic field at the inner radius $B_{\rm in}$ can be estimated as $B_{\rm in}\sim 2 \times 10^{-6}$ G, which is almost the same as that of ISM."975" Then. the dillusion timescale is 5 oy is the Thomson scattering cross section. 7 is the particle velocity normalized by the speed of light and m, is the mass of clectron/ positron."," Then, the diffusion timescale is The synchrotron energy loss of a particle with energy $\varepsilon_e$ is described as where $\sigma_{\rm T}$ is the Thomson scattering cross section, $\beta$ is the particle velocity normalized by the speed of light and $m_e$ is the mass of electron/positron."976"Then. The typical energy loss of the electrons/ positronsAs, with energy 2. can be estimated as This means that the high-energy. electrons/positrons injected into the shocked region lose roughly of the energy by the svnchrotron radiation before diffusing out into ISM."," Then, The typical energy loss of the electrons/positrons $\Delta\varepsilon_e$ with energy $\varepsilon_e$ can be estimated as This means that the high-energy electrons/positrons injected into the shocked region lose roughly of the energy by the synchrotron radiation before diffusing out into ISM."977 Therefore. as in the case of white chwarl pulsars (IXHx11). we can conclude that the radiative energy loss of cleetrons/positrons in the pulsar wind nebula is not so large.," Therefore, as in the case of white dwarf pulsars (KIK11), we can conclude that the radiative energy loss of electrons/positrons in the pulsar wind nebula is not so large."978 The above expressions for the estimate of the energy losses are only applicable to the case that the velocity of a ALSP is subsonic in ISM., The above expressions for the estimate of the energy losses are only applicable to the case that the velocity of a MSP is subsonic in ISM.979 The observed. velocity of MSPs is less than that of canonical pulsars in average sense (Hobbsοἱal.2005)., The observed velocity of MSPs is less than that of canonical pulsars in average sense \citep{Ho05}.980. Llowever. some MSI's have the large velocity ancl a few AISPs actually forms bow shock nebulae 2006).," However, some MSPs have the large velocity and a few MSPs actually forms bow shock nebulae \citep{St03, HB06}."981. In this case. the size of the bow shock is described as (ee... Wilkin 1996) where Vis the velocity of à MSI," In this case, the size of the bow shock is described as (e.g., Wilkin 1996) where $V$ is the velocity of a MSP."982 Due to the assumption of the energy equipartition. the strength of the magnetic fie is The ratio of the Larmor radius of electrons/positrons to the bow shock radius is The fact that rafως is Close to unity supports that clectrons/positrons may escape from the bow shock region.," Due to the assumption of the energy equipartition, the strength of the magnetic field is The ratio of the Larmor radius of electrons/positrons to the bow shock radius is The fact that $r_{\rm g}/R_{\rm bow}$ is close to unity supports that electrons/positrons may escape from the bow shock region."983 Therefore. in the case of &I. high-energy. electrons/positrons can escape with an cllicicney of order unity 2010).," Therefore, in the case of $\kappa\sim 1$, high-energy electrons/positrons can escape with an efficiency of order unity \citep{B08, BA10}."984. Even if we consider the case of &>1. the svnchrotron loss can be estimated by using eqs. (," Even if we consider the case of $\kappa \gg 1$, the synchrotron loss can be estimated by using eqs. ("985L1). (14) and (16) as Therefore. we can conclude again that the radiative energy. loss of clectrons/positrons in the pulsar wind nebula is not so large.,"11), (14) and (16) as Therefore, we can conclude again that the radiative energy loss of electrons/positrons in the pulsar wind nebula is not so large."986" The observed electronpositron spectrum after the propagation in ISM is obtained by solving the cilfusion equation where f(.r2.) is the energy. distribution function of electrons/positrons. D(z.)=Doll|z:/3G0V)? is the diffusion coelficient. (5, is the cooling function of the clectrons/positrons which takes into account svnchrotron emissions and inverse Compton scatterings during the propagation. and ο is the injection term."," The observed electron/positron spectrum after the propagation in ISM is obtained by solving the diffusion equation where $f(t,r,\varepsilon_e)$ is the energy distribution function of electrons/positrons, $D(\varepsilon _e)=D_0(1+\varepsilon_e/3{\rm GeV})^{\delta}$ is the diffusion coefficient, $P(\varepsilon_e)$ is the cooling function of the electrons/positrons which takes into account synchrotron emissions and inverse Compton scatterings during the propagation, and $Q(t,\varepsilon_e,r)$ is the injection term."987 Here we adopt Dui=5.8«1077eni?s 8=1/3. which is consistent with the boron-to-carbon ratio according to the latest GALPROP code.," Here we adopt $D_0=5.8\times 10^{28}{\rm cm}^2{\rm s}^{-1}$, $\delta=1/3$, which is consistent with the boron-to-carbon ratio according to the latest GALPROP code."988 Atovan et al. (, Atoyan et al. (9891995) showed a solution in the case of aninstantaneous injection from a single point-like source. i.e. Οντοι1)2Qu(sJolt ()0(r).,"1995) showed a solution in the case of aninstantaneous injection from a single point-like source, i.e. $Q(t, \varepsilon_e,r) \approx Q_0(\varepsilon_e) \delta(t-t_i)\delta(r)$ ."990" Then the observed spectrum Cf.re,E) would be where 2.9 is the energy of electrons/positrons at the time £C/) and which are cooled down to 2, at the time {, and dauis the ciffusion length given by"," Then the observed spectrum $G(t,r,\varepsilon_e; \tilde{t})$ would be where $\varepsilon_{e,0}$ is the energy of electrons/positrons at the time $\tilde{t} (<t)$ and which are cooled down to $\varepsilon_e$ at the time $t$ , and $d_{\rm diff}$is the diffusion length given by"991.22 we present our dataset aud derive aud compare BH mass values for a sample of NLSIs aud Sis.,2 we present our dataset and derive and compare BH mass values for a sample of NLS1s and S1s.992 Iu 33 we estimate the blue absolute magnitudes aud the stellar velocity dispersious of their bulges., In 3 we estimate the blue absolute magnitudes and the stellar velocity dispersions of their bulges.993 Our results are summarized aud discussed iu 1., Our results are summarized and discussed in 4.994" Firstly. we Lave Isolated a list of 23 NLSI aud 23 S51 galaxies of the northern emisphere from Véron-Cettyetal.(2001) on the basis of their “Siu” aud ""51.07 classificatiou aud of the redshift. % <O.1. chosen to avoid that aand Hines fall in a spectral regione with strouee nieht-sky.e emission lines."," Firstly, we have isolated a list of 23 NLS1 and 23 S1 galaxies of the northern emisphere from \citet{vvg01} on the basis of their “S1n” and “S1.0” classification and of the redshift, z $<0.1$, chosen to avoid that and lines fall in a spectral region with strong night-sky emission lines."995 No other selection criteria were applied., No other selection criteria were applied.996 This sample is complete up to visual maguitude 15.5. corresponcdiug to the 80 per cent of the selected galaxies. aud therefore it is useful for our Qut of this sample. we were able to collect optical spectra for 22 NLS1s and 15 Sis.," This sample is complete up to visual magnitude 15.5, corresponding to the 80 per cent of the selected galaxies, and therefore it is useful for our Out of this sample, we were able to collect optical spectra for 22 NLS1s and 15 S1s."997 In particular 19 NL51s aud 7 S1s were extracted from the public data available in the Isaac Newton Group (ING) Archive., In particular 19 NLS1s and 7 S1s were extracted from the public data available in the Isaac Newton Group (ING) Archive.998 These spectra were obtained for dillerent purposes in 1995. 1996. 1999 and 2000. mostly with the Intermediate Dispersion Spectrograph (IDS) mouuted at the 2.51 Isaac Newtou Teescope (INT. Canary Islauds. Spain). aud the others with the ISIS Double Bean Spectrograph (ISIS) at the L2 Willi Herschel Telescope (WHT. Canary Islauds. Spain).," These spectra were obtained for different purposes in 1995, 1996, 1999 and 2000, mostly with the Intermediate Dispersion Spectrograph (IDS) mounted at the 2.5m Isaac Newton Telescope (INT, Canary Islands, Spain), and the others with the ISIS Double Beam Spectrograph (ISIS) at the 4.2m William Herschel Telescope (WHT, Canary Islands, Spain)."999 Other 3 NLSIs anc 8 Sls were observed directly by us in 2002 September aud in 2003 Janua'y using the Asiago Faint Object Spectrograph aud Camera (AFOSC) mounted at the 1.821 telescope of the Padova Astrotonmidca Observatory (Asiago. It:uly).," Other 3 NLS1s and 8 S1s were observed directly by us in 2002 September and in 2003 January using the Asiago Faint Object Spectrograph and Camera (AFOSC) mounted at the 1.82m telescope of the Padova Astronomical Observatory (Asiago, Italy)."1000 Tables 1 ancl 2 /sumniarize t jeIinstrumental setup and 11e lota wavelength coverage for each observation., Tables \ref{tab1} and \ref{tab2} summarize the instrumental setup and the total wavelength coverage for each observation.1001 All spectra were recuced with the same procedure., All spectra were reduced with the same procedure.1002" The iSla data reduction steps - bias alc flat Ποιά corrections. cosmic rays removal. waveleugtli linearizalon. sky-background. subt'actior aud flux calibration- weο carried out with IRAFpackages""."," The usual data reduction steps - bias and flat field corrections, cosmic rays removal, wavelength linearization, sky-background subtraction and flux calibration- were carried out with IRAF."1003. A one-dimensional spectrum of the nucleus was obtained for each galaxy Ες a number of pixels along the spatial cirection or the basis of the seeing coucditions., A one-dimensional spectrum of the nucleus was obtained for each galaxy summing a number of pixels along the spatial direction on the basis of the seeing conditions.1004 When available. adjacent spectral ranges of the same source were combined together.," When available, adjacent spectral ranges of the same source were combined together."1005 The 1a correction for Galactic extinctior was applied using for each galaxy, Then a correction for Galactic extinction was applied using for each galaxy1006of atomie and molecular absorption lines. all of which are resolved into multiple components or display asymmetric profiles in a manner that indicates more than one contributor.,"of atomic and molecular absorption lines, all of which are resolved into multiple components or display asymmetric profiles in a manner that indicates more than one contributor."1007 These are listed in table 2 and displayed in Fig., These are listed in table 2 and displayed in Fig.1008 1., 1.1009 Neutral hydrogen 21 em line profiles are available for this Galactic position (/jj/2257.2. b;;=+9.6) from the LAB survey (Kalberla et al.," Neutral hydrogen 21 cm line profiles are available for this Galactic position $l_{II}$ =257.2, $b_{II}$ =+9.6) from the LAB survey (Kalberla et al."1010 2005) with a broad component extending from +32 to -52 km s! that spans the radial velocity interval of the atomic and molecular absorption lines., 2005) with a broad component extending from +32 to +52 km $^{-1}$ that spans the radial velocity interval of the atomic and molecular absorption lines.1011" Using the Galactic rotation curve for the 3rd quadrant (Brand Blitz 1993) we find a distance consistent with previous estimates. 4.0 kpe. for 9,2220 km s! based on v4 for CH4300A."," Using the Galactic rotation curve for the 3rd quadrant (Brand Blitz 1993) we find a distance consistent with previous estimates, 4.0 kpc, for $\Theta_0$ =220 km $^{-1}$ based on $_{\rm rad}$ for CH."1012. The center of mass velocity of the binary is still uncertain and may be variable (Uthas et al., The center of mass velocity of the binary is still uncertain and may be variable (Uthas et al.1013 2010) but there may be another way to obtain it., 2010) but there may be another way to obtain it.1014 In the JD 55712 spectrum. the emission lines (e.g. Fe II RMT 42. H Balmer) show two nearly symmetric emission peaks. and the same holds for the weak emergent [O I] 5577. eemission with a separation of about 700 km s! and mean velocity of +6023 km s7! If real. this is corresponds to a distance of 25 kpe.," In the JD 55712 spectrum, the emission lines (e.g. Fe II RMT 42, H Balmer) show two nearly symmetric emission peaks, and the same holds for the weak emergent [O I] 5577, emission with a separation of about 700 km $^{-1}$ and mean velocity of $\pm$ 3 km $^{-1}$ If real, this is corresponds to a distance of $\approx$ 5 kpc."1015 There are no independent interstellar measurements from nearby OB stars in. the archives., There are no independent interstellar measurements from nearby OB stars in the archives.1016 We therefore conclude that the distance to T Pyx is 24 and <5 kpe., We therefore conclude that the distance to T Pyx is $>$ 4 and $\leq$ 5 kpc.1017 The diffuse interstellar bands (DIBs). taken from the list by Friedman et al. (," The diffuse interstellar bands (DIBs), taken from the list by Friedman et al. ("10182010) for those best correlated with extinction. provide an independent estimate of E(B-V) for T Pyx.,"2010) for those best correlated with extinction, provide an independent estimate of E(B-V) for T Pyx."1019 The equivalent widths are given in table 3. the uncertainty is for the measurements.," The equivalent widths are given in table 3, the uncertainty is $\pm$ for the measurements."1020 Combining 5780. 6205. 6196. and 6283 gives E(B-V) = 0.4940.17.," Combining 5780, 6205, 6196, and 6283 gives E(B-V) = $\pm$ 0.17."1021 The strongest DIBs. aand are shown in Fig.," The strongest DIBs, and are shown in Fig."1022 ] along with the interstellar absorption lines listed in Table 2., 1 along with the interstellar absorption lines listed in Table 2.1023 The effect of this revision on the luminosity. both in quiescence and outburst is obvious but the larger extinction also affects the derived properties of the WD based on the continuum from ((Selvelli et al.," The effect of this revision on the luminosity, both in quiescence and outburst is obvious but the larger extinction also affects the derived properties of the WD based on the continuum from (Selvelli et al."1024 2008)., 2008).1025 The increase in E(B-V) increases the UV spectral gradient with wavelength such that the effective temperature of the WD is lower than earlier determinations (but how much this changes depends on a model atmosphere analysis that we will discuss in the next paper in this series)., The increase in E(B-V) increases the UV spectral gradient with wavelength such that the effective temperature of the WD is lower than earlier determinations (but how much this changes depends on a model atmosphere analysis that we will discuss in the next paper in this series).1026 Our first observation. JD 55667. was one day after the discovery and showed a fireball spectrum (see. e.g.Schwarz et al.," Our first observation, JD 55667, was one day after the discovery and showed a fireball spectrum (see, e.g.Schwarz et al."1027 2001) dominated by P Cyent profiles of the H I Balmer lines (Figs., 2001) dominated by P Cygni profiles of the H I Balmer lines (Figs.1028 2 and B.1). He L and and He II. for which the absorption extended without a terminal edge to 22500 km s':D.," 2 and B.1), He I, and and He II, for which the absorption extended without a terminal edge to $\approx$ 2500 km $^{-1}$;."1029 The emission complex at due to C Ul and N UL. was well resolved into the two main contributors without an absorption trough on either.," The emission complex at due to C III and N III, was well resolved into the two main contributors without an absorption trough on either."1030 Subsequent spectra showed only low ionization Fe-peak transitions. Ca Π H and K. and the Na I D doublet. all of which displayed increasingly complex sets of discrete high negative radial velocity absorption components (DACs hereafter).," Subsequent spectra showed only low ionization Fe-peak transitions, Ca II H and K, and the Na I D doublet, all of which displayed increasingly complex sets of discrete high negative radial velocity absorption components (DACs hereafter)."1031 The velocity components identified on the Fe IL and Balmer lines on JD 55703 are at -2964.-1661.-1543.-1433.-1303.-1153.-1056.-952. -799 km s! with mean separation Αν=122+21) km s! and Av<50 km s! for every component based on the Fe II RMT 42lines*:: the minimum width was =5 km s7!.," The velocity components identified on the Fe II and Balmer lines on JD 55703 are at -2964,-1661,-1543,-1433,-1303,-1153,-1056,-952, -799 km $^{-1}$ with mean separation $\Delta v_{rad} =122\pm21$ km $^{-1}$ and $\Delta v \leq 50$ km $^{-1}$ for every component based on the Fe II RMT 42; the minimum width was $\approx 5$ km $^{-1}$."1032 Subsequent velocity shifts are consistent with a linear velocity law for the ejecta., Subsequent velocity shifts are consistent with a linear velocity law for the ejecta.1033 The broader are likely blends of still optically thick individual components., The broader are likely blends of still optically thick individual components.1034 If the width indicates the filament linear thickness then their filling factor may be quite low but their individual optical depths, If the width indicates the filament linear thickness then their filling factor may be quite low but their individual optical depths1035even in the absence of nebulay continua.,even in the absence of nebular continuum.1036 Iu all. the technique of observiug the central stars of PNe with the STIS 50CCD aperture has been very effective aud vields iu accurate V inaguitude for most of the time.," In all, the technique of observing the central stars of PNe with the STIS 50CCD aperture has been very effective and yields an accurate $V$ magnitude for most of the time."1037 In cases where the stars are undetectable. useful limits cau be set ou the stellar fux.," In cases where the stars are undetectable, useful limits can be set on the stellar flux."1038 Iu addition. the high-resolution image (see Shawetal.2006)) provides a clear view of the proximity of field stars from which the viability of ground-based observations of the ceutral star can be judged.," In addition, the high-resolution image (see \citealt{Shaw:06}) ) provides a clear view of the proximity of field stars from which the viability of ground-based observations of the central star can be judged."1039 Very often it is the central stars of bipolar nebulae that are lost iu the continu of the surrounding nebula. which limits our ability to study. the evolution of their progenitors.," Very often it is the central stars of bipolar nebulae that are lost in the continuum of the surrounding nebula, which limits our ability to study the evolution of their progenitors."1040 Iu order to improve siguificantly on our technique for detecting the central star directly. it will be necessary either to obtain high quality spectra and model the nebular coutinmun. or to obtain high-resolution images ofthese PNe in the ultraviolet.," In order to improve significantly on our technique for detecting the central star directly, it will be necessary either to obtain high quality spectra and model the nebular continuum, or to obtain high-resolution images of these PNe in the ultraviolet."1041" As mentioned in the introduction. in the past. most of the ceutral star properties of PNe in the LMC and SAIC have been derived frou, ground-based optical spectroscopic obscrvatious that do not permit a separation of the stellar continui from that of the nebula."," As mentioned in the introduction, in the past, most of the central star properties of PNe in the LMC and SMC have been derived from ground-based optical spectroscopic observations that do not permit a separation of the stellar continuum from that of the nebula."1042 As a result. the central star Iuniuositv aud telperature have to be computed using modeling techniques that rely completely ou nebular 1ieasuremenuts.," As a result, the central star luminosity and temperature have to be computed using modeling techniques that rely completely on nebular measurements."1043 The strongest duplications of this approach is on the stellar luninosities that are determined from the / line flux of the nebula under the assumption that the nebula is optically thick., The strongest implications of this approach is on the stellar luminosities that are determined from the $\beta$ line flux of the nebula under the assumption that the nebula is optically thick.1044 The nebular > line flix is used as alueasurement of the recombination aud therefore the ionizing flux frou the star., The nebular $\beta$ line flux is used as a measurement of the recombination and therefore the ionizing flux from the star.1045 The main problem with this approach of estimating the stellar hunuinositv is that the nebulae are sometimes totally optically thin. or thin in some civections. aud therefore the derived stellar Iunuinosities are lower linüts.," The main problem with this approach of estimating the stellar luminosity is that the nebulae are sometimes totally optically thin, or thin in some directions, and therefore the derived stellar luminosities are lower limits."1046 Iu Figure 2 we show the distribution in the logL|T plane of the central stars analyzed in this oper.," In Figure 2 we show the distribution in the $\rm \log\,L-\log\,T$ plane of the central stars analyzed in this paper."1047 In the cases where the flux at 1686 was not available (MG. 16. Mo 7. Mo 33. SAIP 13. and Mo 17). we have used the II Zaustra temperature to locate the central stars ou the IIR. diagram.," In the cases where the flux at 4686 was not available (MG 16, Mo 7, Mo 33, SMP 43, and Mo 47), we have used the H Zanstra temperature to locate the central stars on the HR diagram."1048 It is well shown that the IT Zaustra temperature uudoerestimutes the temperature of the central stars of optically lun PNe (i.e. Ialer&Jacoby1989:CanemvaldVicmioasx 2000)) but it is accurate for optically thick objects.," It is well known that the H Zanstra temperature underestimates the temperature of the central stars of optically thin PNe (i.e. \citealt{Kj:89, Gv:00}) ) but it is accurate for optically thick objects."1049 We have also used IE Zaustra temperatures for J 39 that has zero measured flux at 1686 aud for SAIP 67 where the measured L686 fux is very unall aud therefore the uncertainty is expected to be very. hieh., We have also used H Zanstra temperatures for J 39 that has zero measured flux at 4686 and for SMP 67 where the measured 4686 flux is very small and therefore the uncertainty is expected to be very high.1050 Figure 3 shows the distribution iu the logLT plane of the total sample of central stars observed wus withZZST in the LMC.," Figure 3 shows the distribution in the $\rm \log\,L-\log\,T$ plane of the total sample of central stars observed by us with in the LMC."1051 The ceutral star parameters obtained in Villaveretal.(2003) fromZZST GO sogra 8271 (Shawctal.2001) for 16 central stars have been added to the 21 ceutral stars shown in Fig., The central star parameters obtained in \cite{Vss:03} from GO program 8271 \citep{Setal:01} for 16 central stars have been added to the 21 central stars shown in Fig.1052 2., 2.1053 We see an identical distribution of the ceutral star locatious iu the WR diagram when we plot the two samples ogether aud in either sample we do not fud auy ceutral star with an effective temperature below 300000 Is. Froiun stellar evolution theory a massive progenitor would evolve very fast along this region of the IIR diagram aud therefore the probability of detection is very small., We see an identical distribution of the central star locations in the HR diagram when we plot the two samples together and in either sample we do not find any central star with an effective temperature below 000 K. From stellar evolution theory a massive progenitor would evolve very fast along this region of the HR diagram and therefore the probability of detection is very small.1054 From uunerical simulations coupled to the central star evolution (Villaverctal.2002a.b).. a very low-inass central star. although it evolves more slowly iu the IIR. diagram. it still needs at least 30000 vyr (after the star leaves the ACB) for a thin PN shell to become ionized aud therefore observable.," From numerical simulations coupled to the central star evolution \citep{Vgm:02,Vmg:02}, a very low-mass central star, although it evolves more slowly in the HR diagram, it still needs at least 000 yr (after the star leaves the AGB) for a thin PN shell to become ionized and therefore observable."1055 Shorter timescales than ~3 0000 yr are needed to observe ionized PNe originating frou intermecdiateanass progenitors (Villaverotal.2002h)., Shorter timescales than $\sim$ 000 yr are needed to observe ionized PNe originating from intermediate-mass progenitors \citep{Vmg:02}.1056. We believe that the coupled evolution between the star and the nebula is reflected in the location of the central stars in the IIR. diagram shown in Fie., We believe that the coupled evolution between the star and the nebula is reflected in the location of the central stars in the HR diagram shown in Fig.1057 3., 3.1058into account (he scattering οἱ electrons by ρου]αννου plasma waves.,into account the scattering of electrons by beam-driven plasma waves.1059 We show. lor the first Gane. that (he generation and absorption of Langninir waves by an electron beam in ihe non-uniform solar corona leads (o the appearance of a break energy in the observed spectrum at the Earth and naturally explains (he observed apparent early injection of low energv electrons.," We show, for the first time, that the generation and absorption of Langmuir waves by an electron beam in the non-uniform solar corona leads to the appearance of a break energy in the observed spectrum at the Earth and naturally explains the observed apparent early injection of low energy electrons."1060 The transport of energetic electrons in the heliospheric plasma is governed by a variety ol different processes (see?.forareview).., The transport of energetic electrons in the heliospheric plasma is governed by a variety of different processes \citep[see][for a review]{Melrose90}.1061 In this work we consider solar energetic electrons propagating along open magnetic field lines and assume their trausport can be described one- ignoring electromagnetic effects., In this work we consider solar energetic electrons propagating along open magnetic field lines and assume their transport can be described one-dimensionally ignoring electromagnetic effects.1062" Under this assumption. the evolution of the electron distribution function. f/(e.c./) [electrons | s] and the spectral energy density of electron. plasma waves ΤΕον) [ergs 7] can be described sell-consistentlv by the following kinetic equations (e.g.?) where Qe/Ok=3e5,fv is the group velocity of Langmuir waves. / is the wavenumber ol a Langmuir wave. ο(0.7)=fue2i is (he plasma wave growth rate. ancl >, and 72, are the collisional and Landau damping rates of waves respectively."," Under this assumption, the evolution of the electron distribution function $f(v,x,t)$ [electrons $^{-4}$ s] and the spectral energy density of electron plasma waves $W(v,x,t)$ [ergs $^{-2}$ ] can be described self-consistently by the following kinetic equations \cite[e.g.][]{Kontar01a}1063 where $\partial \omega/\partial k = 3v_{Te}^2/v$ is the group velocity of Langmuir waves, $k$ is the wavenumber of a Langmuir wave, $\gamma (v,x) =\frac{\pi \omega_{pe}}{n}v^2\frac{\partial f}{\partial v}$ is the plasma wave growth rate, and $\gamma_c$ and $\gamma_L$ are the collisional and Landau damping rates of waves respectively."1064 The first term on the right hand side of both Equation (1)) and (2)) describes the resonant interaction. «y=he. of electrons and Langmuir waves and was first derived by ? and ?..," The first term on the right hand side of both Equation \ref{eqk1}) ) and \ref{eqk2}) ) describes the resonant interaction, $\omega_{pe}=kv$, of electrons and Langmuir waves and was first derived by \citet{Vedenov_etal1962} and \citet{Drummond_Pines1962}."1065 ΤΕ(οαν4) is normalized to the wave energy density fWdh [eres 7] and plays a similar role for plasma waves as the electron distribution Iunction does for particles.," $W(v,x,t)$ is normalized to the wave energy density $\int W dk$ [ergs $^{-3}$ ] and plays a similar role for plasma waves as the electron distribution function does for particles."1066 The initial distribution function is assumed to be a power-law wilh spectral index a in velocily space and has a finite spatial size d at initial time /= 0):, The initial distribution function is assumed to be a power-law with spectral index $\alpha$ in velocity space and has a finite spatial size $d$ at initial time $t=0$ :1067Results from our study. show that the E|A phase is not limited to a small fraction of predmuinately late-tvpe members. but that may massive earlv-tvpes also uudereo this phase.,"Results from our study show that the E+A phase is not limited to a small fraction of predominately late-type members, but that many massive early-types also undergo this phase."1068" To determine if a significant fraction of cluster carly-types have had an E|A phase. s""e consider MS1051 (iom SD. tage~0,5 Gyr. Ho=10 1)."," To determine if a significant fraction of cluster early-types have had an E+A phase, we consider MS1054 $z=0.83$ , $t_{age}\sim 6.5$ Gyr, $H_0=70$ $^{-1}$ )."1069" In this cluster, ~8% of the E-SO members are also E|Αν (see Fig. ??3)."," In this cluster, $\sim8$ of the E-S0 members are also E+A's (see Fig. \ref{morph_hist}) )."1070 To estimate a lower limit ou the uuuber of cluster E-SQOs that have undergone an E|A phase by i—0.5. we asstune the E|A plase can occur at ο Gage~2.6 Cyr) aud is visible or ~1l Covr.," To estimate a lower limit on the number of cluster E-S0's that have undergone an E+A phase by $z\sim0.8$, we assume the E+A phase can occur at $z\lesssim2.5$ $t_{age}\sim 2.6$ Gyr) and is visible for $\sim1$ Gyr."1071 Tn this case. nore than of the E-SQ's in AIS1051 have had an E|A phase. with the restriction fiat we know of uo E|A’s with internal velocity dispersioiS 250|.," In this case, more than of the E-S0's in MS1054 have had an E+A phase, with the restriction that we know of no E+A's with internal velocity dispersions $>250$."1072 However. the true nunber of cary-types that have had an E|A phase can be easilv if we also consider 1) an increasing E|Ao fracjon with redshift: 2) an increasing characteristic E|A ass with redshift: and 3) the couversion of spirals mto early-types.," However, the true number of early-types that have had an E+A phase can be easily if we also consider 1) an increasing E+A fraction with redshift; 2) an increasing characteristic E+A mass with redshift; and 3) the conversion of spirals into early-types."1073 For example. if the fraction of earlv-tvpe E|As increases to at 2OS. the total wmmber of cluster E-SU's that have undergone an E|A phase in MS1051 increases to ~70%.," For example, if the fraction of early-type E+A's increases to at $z>0.8$, the total number of cluster E-S0's that have undergone an E+A phase in MS1054 increases to $\sim70$."1074.. As for the observed restriction that only E|A’s with o<250 exist at 2XOR. an inereasing characteristic mass may mean there is no upper dass iuit on the E|A phenomenon at higher redshifts: even wieltest cluster galaxies (BCC) may have had an E|A shase.," As for the observed restriction that only E+A's with $\sigma<250$ exist at $z\lesssim0.8$, an increasing characteristic mass may mean there is no upper mass limit on the E+A phenomenon at higher redshifts; even brightest cluster galaxies (BCG) may have had an E+A phase."1075 In addition. morphological transformation is likely o plav a prominent role in increasing the earlv-tvpe E|A raction.," In addition, morphological transformation is likely to play a prominent role in increasing the early-type E+A fraction."1076 From comparing CL1358 to MS1051 (Fie. ?7)).," From comparing CL1358 to MS1054 (Fig. \ref{nsigma_hist}) ),"1077 we estimate ~15% of μαςνο (0>200 1)) E-SU's are ransformied from spirals: even if ouly half of hese undergo an E|A phase. the early-type E|A fraction increases by ~20%.," we estimate $\sim45$ of massive $\sigma>200$ ) E-S0's are transformed from spirals; even if only half of these undergo an E+A phase, the early-type E+A fraction increases by $\sim20$."1078.. Any combination of these factors will ouly increase the nuuber of E-SU's that have uncergouc an E|A phase., Any combination of these factors will only increase the number of E-S0's that have undergone an E+A phase.1079 Our estimates on the significance of the E|A phase depend heavily on how loug the E|A phase is visible aud whether E|As evolve into earl-types., Our estimates on the significance of the E+A phase depend heavily on how long the E+A phase is visible and whether E+A's evolve into early-types.1080 However. even our conservative estimate of ~30% establishes the importance of the E|A pliase for à nou-ible number of cluster E-SO’s.," However, even our conservative estimate of $\sim30$ establishes the importance of the E+A phase for a non-neglible number of cluster E-S0's."1081 Tf the true fraction is Jas we demonstrate it can be easily. the E|A phase would be play a critical role in the transformation of spirals iuto the earlv-tvpes that dominate the cluster population.," If the true fraction is, as we demonstrate it can be easily, the E+A phase would be play a critical role in the transformation of spirals into the early-types that dominate the cluster population."1082 To coufiir the importance of the EA phase in the evolution of cluster galaxies. additional observations in lis redshift reginue and higher are needed.," To confirm the importance of the E+A phase in the evolution of cluster galaxies, additional observations in this redshift regime and higher are needed."1083 We suspect hat more massive cluster members underwent an E|A phase at 2>(kN. but whether this also includes the BCC or if there is an upper mass liuit to the E|A phase is unknown.," We suspect that more massive cluster members underwent an E+A phase at $z>0.8$, but whether this also includes the BCG or if there is an upper mass limit to the E+A phase is unknown."1084 For a cluster at :~1.3. 7 find hat the two brightest ineibers show [OITIIJA37277 CLUUISSIOU and chhanced I absorption: the fourth brightest cluster nelber is an E|A. Measuring OT aud Baliner indices of DCCUs at OS—τν<1.3 can establish if DCCUs also undergo an το|A phase. and whether this transition occurs at 2~1.," For a cluster at $z\sim1.3$, \citet{vandokkum:03} find that the two brightest members show $\lambda3727$ emission and enhanced $\delta$ absorption; the fourth brightest cluster member is an E+A. Measuring OII and Balmer indices of BCG's at $0.8<z<1.3$ can establish if BCG's also undergo an E+A phase, and whether this transition occurs at $z\sim1$."1085 Another interesting possibility is to measure mdiees of 21.5 radio galaxies as they cau be the predecessors of lower redshift Μα (e.c.2)..., Another interesting possibility is to measure indices of $z>1.5$ radio galaxies as they can be the predecessors of lower redshift BCG's \citep[e.g.][]{venemans:02}.1086" With observations at higher redshift. we can determine if the E|A phase is related to the bull of star formation. or if it is situply due to ""frostiug by a small fraction of vounecr stars (?).."," With observations at higher redshift, we can determine if the E+A phase is related to the bulk of star formation, or if it is simply due to “frosting” by a small fraction of younger stars \citep{trager:00}."1087 We estimate the ΕΙΔΑ fraction du. intermediate redshift clusters and examine the physical properties of this population using IDIST/WFEDPC?2 iosaics and extensive eround-based spectroscopy., We estimate the E+A fraction in intermediate redshift clusters and examine the physical properties of this population using HST/WFPC2 mosaics and extensive ground-based spectroscopy.1088 Our results are based on spectral types. galaxy colors. magnitudes. IIubble types. and quantitative structural parameters for 500 confirmed members in three clusters (2:=0.33.0.58.&0.83).," Our results are based on spectral types, galaxy colors, magnitudes, Hubble types, and quantitative structural parameters for $\sim500$ confirmed members in three clusters $z=0.33,~0.58,~\&~0.83$ )."1089 We also include measured iuternal velocity dispersions for a subset of 120 ποπο... and estimate dispersions for the rest of the cluster galaxies using the Euudiainoeutal Plane.," We also include measured internal velocity dispersions for a subset of 120 members, and estimate dispersions for the rest of the cluster galaxies using the Fundamental Plane."1090 We find E|A's make up a non-neelieible componcut of the cluster population (~713%)) at intermediate redshifts., We find E+A's make up a non-negligible component of the cluster population $\sim7-13$ ) at intermediate redshifts.1091" They teud to be blucr than the passive members. and we estimate the E|A phase can brightcu a galaxy by as nuceh as AALp,~125 mag."," They tend to be bluer than the passive members, and we estimate the E+A phase can brighten a galaxy by as much as $\Delta M_{Be}\sim1.25$ mag."1092 Although most of them are disk-dominated systems. E|A’s span the rauge in morphological type to include even E-SO members.," Although most of them are disk-dominated systems, E+A's span the range in morphological type to include even E-S0 members."1093 They can be more huumous than £° and can have internal velocity dispersions iu excess of1., They can be more luminous than $L^{\ast}$ and can have internal velocity dispersions in excess of.1094 The majority of these E|As are not associated with merecrs., The majority of these E+A's are not associated with mergers.1095 The diverse nature of E]A members indicates they are drawn from a heterogenous parent population., The diverse nature of E+A members indicates they are drawn from a heterogenous parent population.1096 The characteristics of their descendants can be equally varied: even Inassive earbv-tvpoe members mav have had an E|A phase in their past., The characteristics of their descendants can be equally varied; even massive early-type members may have had an E+A phase in their past.1097 Our study indicates the high dispersion (o>Ww200 1)) ΕΙΑν at 2=0.83 are the logical progenitors of massive carly-types at lower redshift., Our study indicates the high dispersion $\sigma>200$ ) E+A's at $z=0.83$ are the logical progenitors of massive early-types at lower redshift.1098 The cluster E|A’s ave distributed such that the most huninous. hieh dispersion ones are iu our most distant cluster while only £=L.low dispersion E|As exist in our lowest redshift cluster aud Coma (€2?3..," The cluster E+A's are distributed such that the most luminous, high dispersion ones are in our most distant cluster while only $L\lesssim1099L^{\ast}$, low dispersion E+A's exist in our lowest redshift cluster and Coma \citep{caldwell:96,caldwell:99}."1100 The treud of Increasing Iuninositv aud internal velocity dispersion witli redshift provides compelling evidence for mass evolution iu the cluster E|A population., The trend of increasing luminosity and internal velocity dispersion with redshift provides compelling evidence for mass evolution in the cluster E+A population.1101" This galaxy ""dowau-siziug"" is simular to the observed decrease m huninosity of rapidly star-forming field galaxies since +~1 (7).", This galaxy “down-sizing” is similar to the observed decrease in luminosity of rapidly star-forming field galaxies since $z\sim1$ \citep{cowie:96}.1102" Comparison of the E|A luminosity. velocity dispersion. and mass (x 4,07) distributions to the rest of the cluster members indicate this evolution in the E]A iuass distribution is real."," Comparison of the E+A luminosity, velocity dispersion, and mass $\propto r_e\sigma^2$ ) distributions to the rest of the cluster members indicate this evolution in the E+A mass distribution is real."1103 However. we cannot completely rule out that the luminous. massive E|Αν found at higher redshift are due only to au increasing E|A fraction.," However, we cannot completely rule out that the luminous, massive E+A's found at higher redshift are due only to an increasing E+A fraction."1104 Using statistical argunieuts. we find that 230% of cluster E-SUs at :=0.53 have had an E|A phase.," Using statistical arguments, we find that $\gtrsim30$ of cluster E-S0's at $z=0.83$ have had an E+A phase."1105 We cousider this a lower limit as evolution in the E|A fraction auc characteristic mass as well as the conversion of spirals iuto earlv-tvpes can increase the true fraction toA., We consider this a lower limit as evolution in the E+A fraction and characteristic mass as well as the conversion of spirals into early-types can increase the true fraction to.1106. These results show that the E|A phase can play an mportaut role in the transformation of star-forming ealaxies iuto carly-type members., These results show that the E+A phase can play an important role in the transformation of star-forming galaxies into early-type members.1107 Recognizing that our study is based on only three clusters. theseresults can only benefit from. a similar analysis of other clusters within this redshift ranec (0.3<i 0.8).," Recognizing that our study is based on only three clusters, theseresults can only benefit from a similar analysis of other clusters within this redshift range $0.3<z<0.8$ )."1108 Also needed is a comparative study of intermediate 2 field E|As to determine if field aud cluster E|A’s evolve ina similar manucr., Also needed is a comparative study of intermediate $z$ field E+A's to determine if field and cluster E+A's evolve in a similar manner.1109 Previous studies ou the Ποια E|A fraction at 20.2 disagree, Previous studies on the field E+A fraction at $z>0.2$ disagree1110In Figure 1. we display a monochromatic PSF obtained from the CAOS simulations.,In Figure \ref{inputpsf} we display a monochromatic PSF obtained from the CAOS simulations.1111 Even if the SPHERE package of the CAOS system allows simulating different types. of coronagraphs. we preferred to use only a 4-quadrant one for our simulations (Boccalettietal.2008).," Even if the SPHERE package of the CAOS system allows simulating different types of coronagraphs, we preferred to use only a 4-quadrant one for our simulations \citep{Boccaletti08}."1112. In this way. however. our results are still representative às we were not interested in investigating the performances of all the SPHERE instrument coronagraphs.," In this way, however, our results are still representative as we were not interested in investigating the performances of all the SPHERE instrument coronagraphs."1113 The choice of making a simulation with only 64 wavelengths was determined because for more wavelengths. the program saturates our computer memory.," The choice of making a simulation with only 64 wavelengths was determined because for more wavelengths, the program saturates our computer memory."1114 However. CSP requires 269 PSFs at different wavelengths as input and to obtain them. we performed interpolations starting from the ones resulting from the CAOS An early version of the CSP code. described in Bertonet (2006).. has been deeply modified to take variations in the instrument optical design into account.," However, CSP requires 269 PSFs at different wavelengths as input and to obtain them, we performed interpolations starting from the ones resulting from the CAOS An early version of the CSP code, described in \citet{Be06}, has been deeply modified to take variations in the instrument optical design into account."1115 CSP only considers the real part of the image on the lenslet plane and then propagates it through the IFS spectrograph using a Fraunhofer approach. but it can include a treatment of the cross-talk through a parametric approach.," CSP only considers the real part of the image on the lenslet plane and then propagates it through the IFS spectrograph using a Fraunhofer approach, but it can include a treatment of the cross-talk through a parametric approach."1116 The code can be divided into different parts: Where not specified. the simulations were performed assuming a GO spectral type central star with an absolute magnitude of J=3.75 and at a distance of 10 pe from the Sun.," The code can be divided into different parts: Where not specified, the simulations were performed assuming a G0 spectral type central star with an absolute magnitude of J=3.75 and at a distance of 10 pc from the Sun."1117 A total exposure time of 1 hour was generally simulated even if. In some cases. longer exposure times were simulated.," A total exposure time of 1 hour was generally simulated even if, in some cases, longer exposure times were simulated."1118 For these simulations we assumed a readout noise of 10 e—. a dark current of 0.1 e—. and a flat field error of 1077 (hereinafter detector noise).," For these simulations we assumed a readout noise of 10 $e{-}$, a dark current of 0.1 $e{-}$, and a flat field error of $10^{-4}$ (hereinafter detector noise)."1119 High-performance coronagraphs within an extreme AO system like the one adopted in SPHERE. which gets its sampling frequency equal to 20 eycles/pupil. allows imaging of companion objects down to a contrast of 107? within the whole FOV of its IFS (2.5 aresee diagonal). except for separations smaller than -O.1 aresee from the central star.," High-performance coronagraphs within an extreme AO system like the one adopted in SPHERE, which gets its sampling frequency equal to 20 cycles/pupil, allows imaging of companion objects down to a contrast of $10^{-5}$ within the whole FOV of its IFS (2.5 arcsec diagonal), except for separations smaller than $\sim$ 0.1 arcsec from the central star."1120 However. to fulfill the goal of the SPHERE instrument to," However, to fulfill the goal of the SPHERE instrument to"1121to build a theoretical understanding of the galactic magnetic fields (Becketal.1996).,to build a theoretical understanding of the galactic magnetic fields \citep{Beck_etal}.1122. This issue remains unresolved in both the theoretical and numerical senses of the word., This issue remains unresolved in both the theoretical and numerical senses of the word.1123 We note that. because of the large scale separations that have to be captured. the brute-force numerical solution of this problem remains beyond the capacity of currently available computational resources.," We note that, because of the large scale separations that have to be captured, the brute-force numerical solution of this problem remains beyond the capacity of currently available computational resources."1124 On the other hand. the physies of the small-scale magnetic turbulence can be effectively studied as a separate problem and accessed (if only just) by numerical experiment.," On the other hand, the physics of the small-scale magnetic turbulence can be effectively studied as a separate problem and accessed (if only just) by numerical experiment."1125 We hope that by developing a thorough physical understanding of the small-scale magnetic fluctuations. we can approach the problem of their interplay with the large-scale fields and motions.," We hope that by developing a thorough physical understanding of the small-scale magnetic fluctuations, we can approach the problem of their interplay with the large-scale fields and motions."1126 In this work. we concentrate on the structural properties of the small-scale magnetic fields: namely. the intermittency of their spatial distribution and the geometry of the field lines.," In this work, we concentrate on the structural properties of the small-scale magnetic fields: namely, the intermittency of their spatial distribution and the geometry of the field lines."1127 Quantitatively. these are studied in terms of the one-point statistics of the field strength and of the field-line curvature.," Quantitatively, these are studied in terms of the one-point statistics of the field strength and of the field-line curvature."1128" We consider the nonlinear turbulent dynamo to be described by the equations of incompressible MHD: vu- vu- Vp ΕΕ. VB=B +vB... where d/dt=0,+u-V is the convective derivative. (1.X) Is the velocity field. B(r.x) is the magnetic field. and f(r.x) is a random (white in time) large-scale foreing."," We consider the nonlinear turbulent dynamo to be described by the equations of incompressible MHD: = - p + + = +, where $d/dt=\d_t+\vu\cdot\nabla$ is the convective derivative, $\vu(t,\vx)$ is the velocity field, $\vB(t,\vx)$ is the magnetic field, and $\vf(t,\vx)$ is a random (white in time) large-scale forcing."1129 The density p of the plasma is taken to be constant., The density $\rho$ of the plasma is taken to be constant.1130 The incompressibility condition V-u=0Ois. therefore. added to the equations above and serves to determine (or. indeed. to define) pressure.," The incompressibility condition $\nabla\cdot\vu = 0$ is, therefore, added to the equations above and serves to determine (or, indeed, to define) pressure."1131 For the sake ofconvenience. the pressure p and the magnetic field B have been normalized to p and (bz7. respectively.," For the sake ofconvenience, the pressure $p$ and the magnetic field $\vB$ have been normalized to $\rho$ and $(4\pi\rho)^{1/2}$, respectively."1132" The plan of further proceedingsp)"" is as follows.", The plan of further proceedings is as follows.1133 In2.. we review the necessary facts about the kinematic regime of the dynamo.," In, we review the necessary facts about the kinematic regime of the dynamo."1134 These are important. for they form the basis of our understanding of the dynamo and remain surprisingly relevant in the nonlinear regime.," These are important, for they form the basis of our understanding of the dynamo and remain surprisingly relevant in the nonlinear regime."1135 The latter constitutes. our main object of study., The latter constitutes our main object of study.1136 In3.. we describe the results of our numerical experiments and propose a heuristic physical model that explains these results.," In, we describe the results of our numerical experiments and propose a heuristic physical model that explains these results."1137 Conclusions are drawn in4., Conclusions are drawn in.1138. In the weak-field (kinematic) limit. MHD turbulence reduces to the problem of passive advection of a vector’ field by a turbulent velocity field.," In the weak-field (kinematic) limit, MHD turbulence reduces to the problem of passive advection of a vector field by a turbulent velocity field."1139 The magnetic energy grows exponentially and the small-scale folding structure ts formed at the time scale associated with the eddies that turn over the fastest (.e.. in. Kolmogorov turbulence. the viscous-scale eddies).," The magnetic energy grows exponentially and the small-scale folding structure is formed at the time scale associated with the eddies that turn over the fastest (i.e., in Kolmogorov turbulence, the viscous-scale eddies)."1140 Physically. this follows from the fact that the turbulent eddies act on the small-scale fields as a sequence of random linear-shear transformations (see and $2.3)).," Physically, this follows from the fact that the turbulent eddies act on the small-scale fields as a sequence of random linear-shear transformations (see and )."1141 An expanding distribution of the field strength emerges. which is qualitatively explained in terms of the Central Limit Theorem.," An expanding distribution of the field strength emerges, which is qualitatively explained in terms of the Central Limit Theorem."1142 All of the above results can be derived analytically in the framework of the Kraichnan(1968) model of passive advection. which replaces the turbulent velocity with a Gaussian random field à-correlated in time: μυ.," All of the above results can be derived analytically in the framework of the \citet{Kraichnan_ensemble} model of passive advection, which replaces the turbulent velocity with a Gaussian random field $\delta$ -correlated in time: = )."1143 In the context of the small-scale dynamo. this model was first proposed by Kazantsev(1967).," In the context of the small-scale dynamo, this model was first proposed by \citet{Kazantsev}."1144. While the Kazantsev—Kraichnan velocity field is. of course. highly artificial and does not approximate the real turbulent velocity field in any controlled sense. its performance in capturing the essential qualitative and. in some cases. also quantitative. features of the passive advection has been very impressive (on the passive scalar. see recent review by Falkovich.Gawedzki.&Vergas-sola 2001: on the kinematic dynamo. see Kinneyetal.2000.. SCMM02. MCMO2. Schekochihinetal. 2002d)).," While the Kazantsev--Kraichnan velocity field is, of course, highly artificial and does not approximate the real turbulent velocity field in any controlled sense, its performance in capturing the essential qualitative and, in some cases, also quantitative, features of the passive advection has been very impressive (on the passive scalar, see recent review by \citealt{Falkovich_Gawedzki_Vergassola}; on the kinematic dynamo, see \citealt{Kinney_etal}, SCMM02, MCM02, \citealt{SMCM_stokes}) )."1145 This seems to suggest that the statistics of passive advection may be largely universal with respect to the structure of the ambient random flow., This seems to suggest that the statistics of passive advection may be largely universal with respect to the structure of the ambient random flow.1146" In the limit of large Pr,,. the magnetic fluctuations are mostly excited deep in the subviscous range. in which the fluid motions are strongly damped by viscous dissipation."," In the limit of large $\Pr$, the magnetic fluctuations are mostly excited deep in the subviscous range, in which the fluid motions are strongly damped by viscous dissipation."1147 The velocity field “seen” by the magnetic field is. therefore. regular and. in fact. effectively constitutes a flow.," The velocity field “seen” by the magnetic field is, therefore, regular and, in fact, effectively constitutes a ."1148 Most of the relevant statistical results turn out to depend just on the first few, Most of the relevant statistical results turn out to depend just on the first few1149were previously taken into account in [1].. while we will be interested iu the effect of the uext term.,"were previously taken into account in \cite{Kataev:1993be}, while we will be interested in the effect of the next term."1150 It reads The inclusion of the expressions for the two-loop diagrams with massive quark loop insertions. tabulated and taken iuto account in the RuuDec Mathematica package of |15] leads to slight modification of the αἲ- aud a?-correctious to Eq.7y)s' The coustant term of AT(QS) is also affected by the coutributions to the a3-cocticient of Eq.(6) of the diagrams with massive quark loop insertions. evaluated in [16|..," It reads The inclusion of the expressions for the two-loop diagrams with massive quark loop insertions, tabulated and taken into account in the RunDec Mathematica package of \cite{Chetyrkin:2000yt} leads to slight modification of the $a_s^2$ - and $a_s^3$ -corrections to \ref{OS}) ): The constant term of $\Delta\Gamma_3^{(OS)}$ is also affected by the contributions to the $a_s^3$ -coefficient of Eq.(6) of the diagrams with massive quark loop insertions, evaluated in \cite{Bekavac:2007tk}."1151 HToscever. even iu the case of charm-quark loop. these extra ternis are rather small.," However, even in the case of charm-quark loop, these extra terms are rather small."1152 We will neglect these massive-depeudent effects., We will neglect these massive-dependent effects.1153 Fig.1 demonstrate the of the ratio Πλ)=TUL?>bb)πι both in the case of running quark mass and pole quark mass paraimeterisatious., Fig.1 demonstrate the of the ratio $R_b(M_H)=\Gamma(H^0\rightarrow \overline{b}b)/\Gamma_0^{(b)}$ both in the case of running quark mass and pole quark mass parameterisations.1154 The QCD paralucters are fixed as; my=LTGeV and mi(my)=L351GeV [17]. NLO: ACE=253 MeV. and NNLO aud N?LO: A=220 MeV. which correspond to the central values of the results obtained in Ref.," The QCD parameters are fixed as: $m_b=4.7~{\rm GeV}$ and $\overline{m}_b(\overline{m}_b)=11554.34~ {\rm GeV}$ \cite{Penin:2002zv}, NLO: $\Lambda_{\overline{MS}}^{(5)}=253$ MeV and NNLO and $^3$ LO: $\Lambda_{\overline{MS}}^{(5)}=220$ MeV, which correspond to the central values of the results, obtained in Ref."1156 |11] Oue can see. that within the application of the “rumuine” quark mass approach. or as it sometimes called RC-inspired approach. the results do not depeud essentially from the values of the cocfiicicut functions of Eq.(1). calculated," \cite{Kataev:2001kk}1157 One can see, that within the application of the “running” quark mass approach, or as it sometimes called RG-inspired approach, the results do not depend essentially from the values of the coefficient functions of Eq.(1), calculated"1158Classified as genuine primordial binary open clusters although only. a relatively small fraction (~17%)) appear to be able to survive as conspicuous pairs lor more than 25 Mvr.,classified as genuine primordial binary open clusters although only a relatively small fraction $\sim$ ) appear to be able to survive as conspicuous pairs for more than 25 Myr.1159. The distribution of open cluster separations exhibits an apparent peak at 10-15 pc. analogous to the one observed in both LAIC and SAIC.," The distribution of open cluster separations exhibits an apparent peak at 10-15 pc, analogous to the one observed in both LMC and SMC."1160 Here. we have used N-body simulations in an altempt to understand how primordial binary open clusters evolve and what initial orbital elements are required in order (o explain their observed. properties.," Here, we have used $N$ -body simulations in an attempt to understand how primordial binary open clusters evolve and what initial orbital elements are required in order to explain their observed properties."1161 Vazzquez et al. (, Vázzquez et al. (11622010) concluded that available observational evidence indicates that double cluster lifetimes are short and that is what simulations confirm.,2010) concluded that available observational evidence indicates that double cluster lifetimes are short and that is what simulations confirm.1163 Our main conclusions can be summarized as follows., Our main conclusions can be summarized as follows.1164inlerred from X-rays to that. determined from visual extinction (bv dust). it is in principle possible to place constraints on the composition and dust-to-gas ratio of intervening disk malerial.,"inferred from X-rays to that determined from visual extinction (by dust), it is in principle possible to place constraints on the composition and dust-to-gas ratio of intervening disk material."1165 Edge-on sources COUP 419 and COUP 241 are such examples. suggesting (hat the absorbing material is considerably deficient in dust. (an order of magnitude or more relative to to interstellar value).," Edge-on sources COUP 419 and COUP 241 are such examples, suggesting that the absorbing material is considerably deficient in dust (an order of magnitude or more relative to to interstellar value)."1166 Giidelοἱal.(2008). have detected bipolar X-ray jets emanating from the strongly accreting T Taur svstem DG Tau., \citet{Gudel:2008zr} have detected bipolar X-ray jets emanating from the strongly accreting T Tauri system DG Tau.1167 Comparison of the (wo jets indicates (hat one is probably viewed through an excess column of material CVqzz2.7x107! 7) due to the surrounding disk., Comparison of the two jets indicates that one is probably viewed through an excess column of material $N_H\approx2.7\times 10^{21}$ $^{-2}$ ) due to the surrounding disk.1168 Unfortunately. while (is scenario potentially affords a dilferential analvsis of the intervening material. the visual extinction along this line of sight is uncertain. rendering the gas-Lo-dust ratio determination inconclusive.," Unfortunately, while this scenario potentially affords a differential analysis of the intervening material, the visual extinction along this line of sight is uncertain, rendering the gas-to-dust ratio determination inconclusive."1169 Far-ultvaviolet radiation (FUV) is another important source of energy in protoplanetary disks., Far-ultraviolet radiation (FUV) is another important source of energy in protoplanetary disks.1170 The transport of continuum-FUV is largely controlled by dust. ancl as such is highly sensitive (o dust evolution.," The transport of continuum-FUV is largely controlled by dust, and as such is highly sensitive to dust evolution."1171 In particular. as dust settles towards the disk midplane the upper parts of the disk become increasingly. transparent to EUV photons. reducing the FUV opacily to. X-rays levels.," In particular, as dust settles towards the disk midplane the upper parts of the disk become increasingly transparent to FUV photons, reducing the FUV opacity to X-rays levels."1172 In contrast. the propagation of N-ravs is impeded almost entirely bv gas after only a modest degree of dust evolution (e~ 0.1).," In contrast, the propagation of X-rays is impeded almost entirely by gas after only a modest degree of dust evolution $\epsilon\sim0.1$ )."1173 The deposition of N-rays with energies (vpical of the steady coronal emission from T. Tauri stus (E~1— 2keV) will be deposited in the upper lavers of the disk - the so-called warm molecular laver., The deposition of X-rays with energies typical of the steady coronal emission from T Tauri stars $E\sim1-2$ keV) will be deposited in the upper layers of the disk - the so-called warm molecular layer.1174 In the models ol Aikawa&Nomura(2006) the warm laver of gaseous CO al R=200AU has a typical vertical column density of Ny(CO)~I0Pem >7., In the models of \citet{Aikawa:2006uq} the warm layer of gaseous CO at R=200AU has a typical vertical column density of $N_Z$ $\sim10^{18}$ $^{-2}$.1175 We expect this laver to be verv optically thick. τοι07100. when viewed at (he small grazing angles (75 degrees) appropriate for impinging stellar X-rays.," We expect this layer to be very optically thick, $\tau_{1keV}\sim100$, when viewed at the small grazing angles $\sim5$ degrees) appropriate for impinging stellar X-rays."1176 The impenetrability of this laver is mitigated somewhat at higher energies (IZο” IOkeV) where the photoelectric cross-sections becomes small compared to the scattering cross-section., The impenetrability of this layer is mitigated somewhat at higher energies $E\sim10$ keV) where the photoelectric cross-sections becomes small compared to the scattering cross-section.1177 In this ease. Thomson scattering by gas can be an effective means," In this case, Thomson scattering by gas can be an effective means"1178overshoot laver ε becomes much greater (han unity.,overshoot layer $\epsilon$ becomes much greater than unity.1179 Therefore. the solution of equation (1) for the two extreme cases €<1 and e>1 covers both parts of the tachocline.," Therefore, the solution of equation (1) for the two extreme cases $\epsilon \ll 1$ and $\epsilon \gg 1$ covers both parts of the tachocline."1180 Here we use a magnetic field profile twpical of the Sun. namely which has opposite signs in (he northern and southern hemispheres 1997).," Here we use a magnetic field profile typical of the Sun, namely which has opposite signs in the northern and southern hemispheres \citep{gil97}."1181. The magnetic field profile (2) leads to a=agp. where ag=D4/(2048/1xp).," The magnetic field profile (2) leads to $\alpha=\alpha_0 \mu$, where $\alpha_0=B_0/(2\Omega_0 R_0 \sqrt{4\pi1182\rho})$."1183 In the next subsections we will give analvtical solutions aud dispersion relations of global shallow water MIID waves lor each case separately., In the next subsections we will give analytical solutions and dispersion relations of global shallow water MHD waves for each case separately.1184 Tteduced eravitv in the strongly stable radiative part of the tachocline can be estimated as 5001.5 - 10! emnes. 7? (Schecteretal.2001)., Reduced gravity in the strongly stable radiative part of the tachocline can be estimated as 500–1.5 $\cdot$ $^4$ $\cdot$ $^{-2}$ \citep{sch01}.1185". Then. for the laver thickness //,—10? cm we gel e= 45 - .E 0.13 (where 0422.6 5 !.. Ry=5- 10! em have been used)."," Then, for the layer thickness $H_0$ $^9$ cm we get $\epsilon=$ 4.5 $\cdot$ $^{-3}$ – 0.13 (where $\Omega_0$ =2.6 $\cdot$ $^{-6}$ $^{-1}$, $R_0$ $\cdot$ $^{10}$ cm have been used)."1186 Therefore. e<1 is a good approximation in the radiative part of the tachocline.," Therefore, $\epsilon \ll 1$ is a good approximation in the radiative part of the tachocline."1187" In this approximation and using the weak magnetic field limit. ie. a2< 1. equation (1) leads to the spheroidal wave equation (second order terms with ε and a7 are neglected). whose typical solutions are the spheroidal wave functions &,, 1964).. where n plavs (he role of the poloidal wavenumber."," In this approximation and using the weak magnetic field limit, i.e. $\alpha_0^2 \ll 1$ , equation (1) leads to the spheroidal wave equation (second order terms with $\epsilon$ and $\alpha_0^2$ are neglected), whose typical solutions are the spheroidal wave functions $S_{sn}$ \citep{abr64}, where $n$ plays the role of the poloidal wavenumber."1188 The approximate solution can be written aswhere Wy is the amplitude aud, The approximate solution can be written aswhere $u_0$ is the amplitude and1189found.,found.1190 Instead. an 10/20 either-or domain ts present.," Instead, an 1O/2O either-or domain is present."1191 Although the presence of the resorance may be not sufficient to excite the beat pulsation and other factors may be necessary (like for the F+10 models discussec in the previous section). it is hard to identify these factors ard prove their necessity.," Although the presence of the resonance may be not sufficient to excite the beat pulsation and other factors may be necessary (like for the F+1O models discussed in the previous section), it is hard to identify these factors and prove their necessity."1192 Therefore. we do not presume to know which mechanism. resonant or non-resonant. underlies the ¢ouble-overtone behaviour we found.," Therefore, we do not presume to know which mechanism, resonant or non-resonant, underlies the double-overtone behaviour we found."1193 Although some double-overtone models were found. the overall results are not satisfactory.," Although some double-overtone models were found, the overall results are not satisfactory."1194 The computed double-overtone comains are narrow and are located at Pj.s0.3d. Le. at the short-period end of the observed domain of the double-overtone pulsation in the LMC.," The computed double-overtone domains are narrow and are located at $P_1\approx 0.3\,{\rm d}$, i.e. at the short-period end of the observed domain of the double-overtone pulsation in the LMC."1195 No double-overtone models were fould at longer periods. particularly in a period range. O5d«P<0.9d. where the double-overtone pulsation is the most comnon form of pulsation in the LMC.," No double-overtone models were found at longer periods, particularly in a period range, $0.5\,{\rm d}<P_1<0.9\,{\rm d}$, where the double-overtone pulsation is the most common form of pulsation in the LMC."1196 Surprisingly. in many models only pulsation in the second overtone ts possible.," Surprisingly, in many models only pulsation in the second overtone is possible."1197 Observationally. this form of pulsation Is very rare. as only 14 such objects are identified in the LMC (Soszynsski et al. 2008..," Observationally, this form of pulsation is very rare, as only 14 such objects are identified in the LMC (Soszyńsski et al. \cite{so08b},"1198 still. observational selection can be a factor here).," still, observational selection can be a factor here)."1199 Therefore. also for the 10+20 double-overtone Cepheids the convective hydrocode fails to reproduce the observed modal selection.," Therefore, also for the 1O+2O double-overtone Cepheids the convective hydrocode fails to reproduce the observed modal selection."1200 This conclusion is not as strict as for the F+!lO Cepheids. because only one set of convective parameters was explored.," This conclusion is not as strict as for the F+1O Cepheids, because only one set of convective parameters was explored."

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