ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target22008) measurecl rotational velocities of he primary stars in the classical Aleols.,2008) measured rotational velocities of the primary stars in the classical Algols.3 At least seven systems out of 23 binaries rotate faster which cüller from svnchronize rotation by more han a [actor of two., At least seven systems out of 23 binaries rotate faster which differ from synchronize rotation by more than a factor of two.4 Egeleton (2000) discussed. evolution of Aleol type systems and concluded that they could not o* evolved. to their present. status without having lost substantial mass and angular momentum., Eggleton (2000) discussed evolution of Algol type systems and concluded that they could not be evolved to their present status without having lost substantial mass and angular momentum.5Phe fundamental xwameters of well-observed. detached: ancl semi-detached Algols are compiled and analvsed by. Ibanogllu et al. (,The fundamental parameters of well-observed detached and semi-detached Algols are compiled and analysed by Ibanoğllu et al. (62006) ο reveal some possible implications for their nuclear and angular momentum evolution.,2006) to reveal some possible implications for their nuclear and angular momentum evolution.7 They arrived at a result that he mass-ratio of detached. Algols is larger than unity., They arrived at a result that the mass-ratio of detached Algols is larger than unity.8 As he system evolved-olf from the main-sequence mass transfer, As the system evolved-off from the main-sequence mass transfer9for all Z.,for all $Z$.10 This equation means that the absolute value of D4(Z) should be chosen to be at most of order of 1/E? for all Z., This equation means that the absolute value of $D_{1}(Z)$ should be chosen to be at most of order of $1/E^{2}$ for all $Z$ .11" In fact. if |[D4|291/E? for some Z. we can easily see that equation (53)) breaks down in the range L/E?«i?<|D,)."," In fact, if $|D_{1}|\gg 1/E^{2}$ for some $Z$, we can easily see that equation \ref{sol3}) ) breaks down in the range $1/E^{2}\ll \y^2 \ll12 |D_{1}|$ ."13" Hence. we can discuss the increase of AM?~f/f, iv the range 7>L/P under the choice of D4=0. namely. according to the following equation It is now clear that we have M?~£P in therange l/E«at<1. and the outflows can pass through the fast-inagnetosonic point (i.e.. AP=1/E?) located on the radius RO,cEV,"," Hence, we can discuss the increase of $\mm^{2}\simeq14 E_{k}/E_{m}$ in the range $\y \gg 1/E$ under the choice of $D_{1}=0$, namely, according to the following equation It is now clear that we have $\mm^{2}\simeq \y$ in therange $1/E \ll \y \ll 1$, and the outflows can pass through the fast-magnetosonic point (i.e., $\mm^{2}=1/E^{2/3}$ ) located on the radius $R\Omega_{F}\simeq E^{1/3}$."15" D the outflows can arrive at the radius £1.4. the equipartition Ej=E,, between magnetic and kinetic energies is realized."," If the outflows can arrive at the radius $\y \simeq 1.4$, the equipartition $E_{k}=E_{m}$ between magnetic and kinetic energies is realized."16 In the asymptotic region à1 we can confirm the logarithmic increase of AP given by APsIn(272/AL) toward the complete conversion of magnetic to kinetic enerev., In the asymptotic region $\y \gg 1$ we can confirm the logarithmic increase of $\mm^{2}$ given by $\mm^{2}\simeq\ln(2\y^{2}/\mm^{2})$ toward the complete conversion of magnetic to kinetic energy.17" The numerical solution of equation (54)) is shown in Figure baa. The corresponding poloidal velocity αρ) and Lorentz [actor 5C75) are also shown in Figure 5bb. where the Lorentz factor is given by 5=E—u,/(£M7)."," The numerical solution of equation \ref{model}) ) is shown in Figure \ref{fig:gs-mm}a a. The corresponding poloidal velocity $u_p(\y)$ and Lorentz factor $\gamma(\y)$ are also shown in Figure \ref{fig:gs-mm}b b, where the Lorentz factor is given by $\gamma = E - u_p/(\xi \mm^2)$."18 In this case. we obtain ελ<Il anvwhere forthe ratio of the poloidal electric to toroidal electric magnetic field amplitude (see Fig.5cec).," In this case, we obtain $\xi(\y)<1$ anywhere forthe ratio of the poloidal electric to toroidal electric magnetic field amplitude (see \ref{fig:gs-mm}c c)."19 Note that the Lorentz factor includes both the poloidal and toroidal motion., Note that the Lorentz factor includes both the poloidal and toroidal motion.20" In Fig.5bb. the difference between 4,-value and 5-value near the light. evlinder is due to the dominated toroidal motion of the plasma. ("," In \ref{fig:gs-mm}b b, the difference between $u_p$ -value and $\gamma$ -value near the light cylinder is due to the dominated toroidal motion of the plasma. ("21Just around the light evlinder ££~7). the properties of the flow shown in Fie.,"Just around the light cylinder $\y \sim \y_{\rm L}$, the properties of the flow shown in Fig."22 5 may be incorrect because of our approximations. but we can expect correct features al least near and outside the fast magnetosonic point. > ty. )," \ref{fig:gs-mm} may be incorrect because of our approximations, but we can expect correct features at least near and outside the fast magnetosonic point, $\y \geq \y_{\rm F}$ . )"23 Next. we discuss the field configuration given by equation (52)).," Next, we discuss the field configuration given by equation \ref{sol2}) )."24 Thejet flows may be confined by an external pressure (see. e.g.. Li 1993: Degelman Li 1994: Fendt 1997).," Thejet flows may be confined by an external pressure (see, e.g., Li 1993; Begelman Li 1994; Fendt 1997)."25" If the shape ZQ),=ICE) of the last (lux surface Y=Wp is determined by the outer boundary condition. the fnuction o(Z) in equation (52)) is fixed."," If the shape $Z\Omega_{F}=H(\y)$ of the last flux surface $\Psi=\Psi_{0}$ is determined by the outer boundary condition, the function $D_{2}(Z)$ in equation \ref{sol2}) ) is fixed."26" For example. let us consider the racial last {lux surface al an angle 2=0, wilh the pole axis direction. ("," For example, let us consider the radial last flux surface at an angle $R/Z \equiv \theta_{0}$ with the pole axis direction. ("27"Because the solution can be applied only (ο jets with small opening angles. we must require thal 0, is al mostof order of 1/ E.)","Because the solution can be applied only to jets with small opening angles, we must require that $\theta_{0}$ is at mostof order of $1/E$ .)"28 Using the function AP=a) derived [rom equation (54)). we obtain alonga [lux ποπο.VOR.Z)= constant. where 7= ZOL/E.," Using the function $\mm^2=\mm^2(\y)$ derived from equation \ref{model}) ), we obtain alonga flux function$\Psi(R,Z)=$ constant, where $\z \equiv29 Z\Omega_F/E$ ."30" Note that lor the value of M?(8,2) in equation (55)) we use the [unctionM?= AM?(P). where the variable ? should be replaced to 052."," Note that for the value of $\mm^2(\theta_0 \z)$ in equation \ref{eq:d_Z}) ) we use the function$\mm^2=\mm^2(\y)$ , where the variable $\y$ should be replaced to $\theta_0 \z$ ."31" Then. in the range 70,/2« lof Z. equations (52)) ancl (55)) leads"," Then, in the range $\theta_{0}Z\Omega_{F}/E \ll 1$ of $Z$ , equations \ref{sol2}) ) and \ref{eq:d_Z}) ) leads"32other hand. a measured anisotropy. larger than the 2356 we obtained for D0656--14 would be. from and(43).. an indication of the existence of an even vounger. nearby source. aa longer lifetime { of the PWN.,"other hand, a measured anisotropy larger than the $\approx$ we obtained for B0656+14 would be, from and, an indication of the existence of an even younger, nearby source, a longer lifetime $T$ of the PWN."33 We also note that the predicted Πας from PSR. D0656--14 appears (o overpredict the observed flux above GGeV. This implies more severe constraints on the pulsar output. whereas Geminga's parameters are not that severely constrained by CR. positron observations.," We also note that the predicted flux from PSR B0656+14 appears to overpredict the observed flux above GeV. This implies more severe constraints on the pulsar output, whereas Geminga's parameters are not that severely constrained by CR positron observations."34 We remark that Galactic CR. including electrons and. positrons. are subjected to solar modulation at energies below 210 GGeV. The encounter of (hese particles with the solar wind and imbedcded magnetic field causes a heliospheric anisotropy. (hat is primarily determined bv the combined modulation effects of convection. diffusion. and drilts - all solar cevcle dependent.," We remark that Galactic CR, including electrons and positrons, are subjected to solar modulation at energies below $\approx$ GeV. The encounter of these particles with the solar wind and imbedded magnetic field causes a heliospheric anisotropy that is primarily determined by the combined modulation effects of convection, diffusion, and drifts - all solar cycle dependent."35 Drifts will cause (his anisotropy to have a 22-vear cvele., Drifts will cause this anisotropy to have a 22-year cycle.36 CR electrons. ancl positrons al GGeV. at Earth may. therelore exhibit a heliospheric anisotropy. of up to a few percent. assuming (hat thev enter the heliosphere isotropically 2004)..," CR electrons and positrons at GeV at Earth may therefore exhibit a heliospheric anisotropy of up to a few percent, assuming that they enter the heliosphere isotropically \citep{2004ApJ...602..993P}."37 It will be an interesting exercise to determine how this anisotropy will change if the LIS is anisotropic., It will be an interesting exercise to determine how this anisotropy will change if the LIS is anisotropic.38 However. the anisotropy that we predict here is the largest above GGeV. an energv range αἱ which only the PAMELA mission (Doezioetal.2004) mav be," However, the anisotropy that we predict here is the largest above GeV, an energy range at which only the PAMELA mission \citep{2004NuPhS.134...39B} may be"39"of the background to the total flux from the stars is low for both observations, from to depending on the image.","of the background to the total flux from the stars is low for both observations, from to depending on the image."40" Therefore, photometric errors are not dominated by fluctuations in the background."," Therefore, photometric errors are not dominated by fluctuations in the background."41" We used a sliding median filter to select and trim outliers in flux and position greater than 50, representing and of the data for the first and second visits, respectively."," We used a sliding median filter to select and trim outliers in flux and position greater than $\sigma$, representing and of the data for the first and second visits, respectively."42" We also discarded the first half-hour's worth of observations, which is affected by significant telescope jitter before stabilization."," We also discarded the first half-hour's worth of observations, which is affected by significant telescope jitter before stabilization."43 The final number of photometric measurements used is 7277 and 7362., The final number of photometric measurements used is 7277 and 7362.44 The raw time series are presented in the top panel of Figure 1.., The raw time series are presented in the top panel of Figure \ref{fig:spitzerlightcurves}.45" We find that the point-to-point scatter in the photometry gives a typical signal-to-noise ratio (S/N) of 280 per image, which corresponds to of the theoretical signal-to-noise."," We find that the point-to-point scatter in the photometry gives a typical signal-to-noise ratio (S/N) of 280 per image, which corresponds to of the theoretical signal-to-noise."46" T'herefore, the noise is dominated by Poisson statistics."," Therefore, the noise is dominated by Poisson statistics."47" In order to determine the transit parameters and associated uncertainties from the ttime series we used a transit light curve model multiplied by instrumental decorrelation functions, as described by Désertetal.(2011a)."," In order to determine the transit parameters and associated uncertainties from the time series we used a transit light curve model multiplied by instrumental decorrelation functions, as described by \cite{desert11a}."48. The transit light curves were computed with the IDL transit routine from Mandel&Agol(2002)., The transit light curves were computed with the IDL transit routine from \cite{mandel02}.49" For the present case we allowed for a single free parameter in the model, which is the planet-to-star radius ratio equivalently, the depth, in the absence of limb R,/R,darkening)."," For the present case we allowed for a single free parameter in the model, which is the planet-to-star radius ratio $R_p/R_\star$ (or equivalently, the depth, in the absence of limb darkening)."50"(or The normalized orbital semi-major axis (system scale) a/R,, the impact parameter b, the period P, and the time of mid transit T, were held fixed at the values derived from the light curve, as reported by Batalhaetal.(2011) and summarized below in Sect. 7.."," The normalized orbital semi-major axis (system scale) $a/R_\star$, the impact parameter $b$, the period $P$, and the time of mid transit $T_c$ were held fixed at the values derived from the light curve, as reported by \cite{Batalha:11} and summarized below in Sect. \ref{sec:discussion}."51" Limb darkening is small atmicron,, but was nevertheless included in our modeling using the 4-parameter law by Claret(2000) and theoretical coefficients published by Sing(2010)."," Limb darkening is small at, but was nevertheless included in our modeling using the 4-parameter law by \cite{Claret:00} and theoretical coefficients published by \cite{Sing:10}."52". The Spitzer//IRAC photometry is known to be systematically affected by the so-called “pixel-phase effect"" (see,e.g.,Charbonneauetal.2005;Knutsonetal.2008)."," The /IRAC photometry is known to be systematically affected by the so-called “pixel-phase effect” \citep[see,53e.g.,][]{charbonneau05,knutson08}."54. This effect is seen as oscillations in the measured fluxes with a period corresponding to that of the telescope pointing jitter., This effect is seen as oscillations in the measured fluxes with a period corresponding to that of the telescope pointing jitter.55" For the first visit this period was 70 min, and the amplitude of the oscillations was approximately peak-to-peak; for the second visit the period was 35 min, and the amplitude about196.."," For the first visit this period was 70 min, and the amplitude of the oscillations was approximately peak-to-peak; for the second visit the period was 35 min, and the amplitude about."56 We decorrelated our signal in each channel using a linear function of time for the baseline (two parameters) and a quadratic function of the PSF position (four parameters) to correct the data for each channel., We decorrelated our signal in each channel using a linear function of time for the baseline (two parameters) and a quadratic function of the PSF position (four parameters) to correct the data for each channel.57 We performed a simultaneous Levenberg-Marquardt least-squares fit to the data (Markwardt2009) to determine the transit and instrumental model parameters (7 in total).," We performed a simultaneous Levenberg-Marquardt least-squares fit to the data \citep{markwardt09}58 to determine the transit and instrumental model parameters (7 in total)."59" The errors on each photometric point were assumed to be identical, and were set to the rms residual of the initial best fit."," The errors on each photometric point were assumed to be identical, and were set to the rms residual of the initial best fit."60" To obtain an estimate of the correlated and systematic errors in our measurements etal.2006) we used the residual permutation bootstrap(Pont technique, or “Prayer Bead” method, as described by Désertetal.(2009)."," To obtain an estimate of the correlated and systematic errors in our measurements \citep{pont06} we used the residual permutation bootstrap technique, or “Prayer Bead” method, as described by \citet{desert09}."61". In this method the residuals of the initial fit are shifted systematically and sequentially by one frame, and then added to the transit light curve model before fitting again."," In this method the residuals of the initial fit are shifted systematically and sequentially by one frame, and then added to the transit light curve model before fitting again."62" We considered asymmetric error bars spanning of the points above and below the median of the distributions to derive the lo uncertainties for each parameter, as described by Désertetal.(2011b)."," We considered asymmetric error bars spanning of the points above and below the median of the distributions to derive the $1\sigma$ uncertainties for each parameter, as described by \citet{desert11b}."63". 'The bottom panel of Figure 1 shows the best-fit model superimposed on the observations from the two visits combined, with the data binned in 36mmin bins for clarity (295 points per bin)."," The bottom panel of Figure \ref{fig:spitzerlightcurves} shows the best-fit model superimposed on the observations from the two visits combined, with the data binned in min bins for clarity (295 points per bin)."64" The transit depths at ((after removing limb-darkening effects) are 353*112 ppm for the first visit and 339777, for the second, which are in good agreement with each other."," The transit depths at (after removing limb-darkening effects) are $353^{+115}_{-133}$ ppm for the first visit and $339^{+85}_{-110}$ for the second, which are in good agreement with each other."65" The weighted average depth of 344+85 is consistent with the non-limb-darkened value of 376+9 ppm derived from the light curve (Batalhaetal. well within the lo errors, strongly suggesting the 2011)transit is achromatic, as expected for a planet."," The weighted average depth of $344 \pm 85$ is consistent with the non-limb-darkened value of $376 \pm 9$ ppm derived from the light curve \citep{Batalha:11} well within the $\sigma$ errors, strongly suggesting the transit is achromatic, as expected for a planet."66 'The above oobservations provide a useful constraint on the kinds of false positives (blends) that may be mimicking the ssignal., The above observations provide a useful constraint on the kinds of false positives (blends) that may be mimicking the signal.67" For example, if Kepler-10 were blended with a faint unresolved background eclipsing binary of much later spectral type that manages to reproduce the transit depth in the ppassband, the predicted depth at mmay be expected to be larger because of the higher flux of the contaminating binary at longer wavelengths compared to Kepler-10."," For example, if Kepler-10 were blended with a faint unresolved background eclipsing binary of much later spectral type that manages to reproduce the transit depth in the passband, the predicted depth at may be expected to be larger because of the higher flux of the contaminating binary at longer wavelengths compared to Kepler-10."68" Since the transit depth we measure in the near infrared is about the same as in the optical, this argues against blends composed of stars"," Since the transit depth we measure in the near infrared is about the same as in the optical, this argues against blends composed of stars"69erowtli coutinues as dark matter halo mergers proceed to low redshifts.,growth continues as dark matter halo mergers proceed to low redshifts.70 Gas accretion is thought to play a critical role in fucling the early stages of black hole erowth (???).. and this may explain the tightness of the AIpy0 relation (22777)..," Gas accretion is thought to play a critical role in fueling the early stages of black hole growth \citep{david:1987:eag, 71kauffmann:2000:ume, merloni:2004:ags}, and this may explain the tightness of the $_{\rm BH}-\sigma$ relation \citep{Burkert:2001:smbh,72haehnelt:2000:cbh, Dimatteo:2005, Kaz:2005:msigma, 73Robertson:2006:msigma}."74 Since high redshift galaxies are thought to be especially eas-rich. each merecr brings a fresh supply of gas to the ceuter of the galaxy. aud new fuel to the erowine supermassive black hole (??)..," Since high redshift galaxies are thought to be especially gas-rich, each merger brings a fresh supply of gas to the center of the galaxy, and new fuel to the growing supermassive black hole \citep{Mihos:1994:gasmerger, Dimatteo:2003:bhgrowth}."75 From a combination of gas accretion and binary black hole coalesceuce. it is thought that these Pop IIT-seuerated seeds may form the SMDITS we observe today (?)..," From a combination of gas accretion and binary black hole coalescence, it is thought that these Pop III-generated seeds may form the SMBHs we observe today \citep{Soltan:1982:smbh, Schneider:02topheavy}."76 During a galaxw merger. cach black hole sinks to the center of the new ealaxy potential due to ανασα] friction aud eveutually becomes bound as a binary (??)..," During a galaxy merger, each black hole sinks to the center of the new galaxy potential due to dynamical friction and eventually becomes bound as a binary \citep{Kaz:2005:msigma, Escala:2005:gas}."77 Dynamical friction then coutiuues to shriuk the orbit until the binary is lard (ie. the separation between each black hole. appr. is such that the svstem tends to lose energv curing stellar encounters) (?)..," Dynamical friction then continues to shrink the orbit until the binary is hard (i.e, the separation between each black hole, $_{\rm BBH}$, is such that the system tends to lose energy during stellar encounters) \citep{heggie:07}."78 Thereatter. further decay is mediated by 3-body scattering with the züubieut stellar backeround uutil the binary becomes so close that the orbit can lose energv via eravitational radiation.," Thereafter, further decay is mediated by 3-body scattering with the ambient stellar background until the binary becomes so close that the orbit can lose energy via gravitational radiation."79 Iu studies of static. spherical potentials. it nav be difficult for stellar eucouuters aloue to cause the αν to transition between the 3-body scattering phase and the gravitational radiation reguue (7).," In studies of static, spherical potentials, it may be difficult for stellar encounters alone to cause the binary to transition between the 3-body scattering phase and the gravitational radiation regime \citep{milos:2003:lem}."80 However. in gasrich or non-splherical svstenis. the binary rapidly judeus and coalesces iuto one black hole. οιτας copious gravitational radiation iu the process (???77)..," However, in gas-rich or non-spherical systems, the binary rapidly hardens and coalesces into one black hole, emitting copious gravitational radiation in the process \citep{Mayer:2007:smbh, Kaz:2005:msigma, berczik:2006:emb, sigurdsson:03, KHB:2006fl}."81" Iu our previous work. we calculated the cosinological ucrger rate for black holes between 200 - 3&10*AL, roni redshift 19-0 022.(hereafter.MOT)).."," In our previous work, we calculated the cosmological merger rate for black holes between 200 - $3 \times 10^7 \Msun$ from redshift 49-0 \citep[][(hereafter, M07)]{KHB:2007bhgrowth, KHB:2008bhgrowth}."82 Our approach combined high-resolution. snall-volume cosmological N-xlv smuulations with analytic prescriptions for the dynamics of nereius black holes below our resolution iuit: this allowed us to explore differeut black hole erowtl mechanisms and seed formation scenarios while also accurately simulating the rich aud varied merecr ustory of the host dark matter halos.," Our approach combined high-resolution, small-volume cosmological N-body simulations with analytic prescriptions for the dynamics of merging black holes below our resolution limit; this allowed us to explore different black hole growth mechanisms and seed formation scenarios while also accurately simulating the rich and varied merger history of the host dark matter halos."83 Iu this paper. we calculate the eravitational wave sienal from the black hole mergers involved in asseiibling a supermassive black hole at the center of the Mille Wav analogue our simulation volume.," In this paper, we calculate the gravitational wave signal from the black hole mergers involved in assembling a supermassive black hole at the center of the Milky Way analogue our simulation volume."84 The volume is designed to provide one possible evolutionary path for a region like our Local Group. aud as such. it should contain supernassive black holes ou the light cud of the supermassive black hole mass spectrum the sweet spot oe1 SMIBIT mass for LISA observations.," The volume is designed to provide one possible evolutionary path for a region like our Local Group, and as such, it should contain supermassive black holes on the light end of the supermassive black hole mass spectrum – the sweet spot in SMBH mass for LISA observations."85 We inchide all 1ο niereers that have occurred from redshift 19 to the xeseut epoch within a 1000 Ape? volune of the Universe iat represeuts a Local Croup type of environment., We include all the mergers that have occurred from redshift 49 to the present epoch within a 1000 $^3$ volume of the Universe that represents a Local Group type of environment.86 We found that the exavitational wave sources revealed oei this volume are from much ligher mass ratio merecrs iat the mergers predicted to be involved iu assenibliug je luost massive SMDIIS., We found that the gravitational wave sources revealed in this volume are from much higher mass ratio mergers that the mergers predicted to be involved in assembling the most massive SMBHs.87" Most of the LISA) scicuce and data analysis community las been auticipating nore equal mass mergers. aud have developed extcusive eravitational wave templates aud paramcter extraction echuiques based on the asswuption that the black hole ünaries are order unitv mass ratio svstenis (7.οι),"," Most of the LISA science and data analysis community has been anticipating more equal mass mergers, and have developed extensive gravitational wave templates and parameter extraction techniques based on the assumption that the black hole binaries are order unity mass ratio systems \citep[][e.g.]{Babak:2008}."88 While this may be true for the most massive SAIBIIs. we find that the black holes in our volume experieuce uergeers with mass ratios as high as Mo/M4=10000: l.," While this may be true for the most massive SMBHs, we find that the black holes in our volume experience mergers with mass ratios as high as ${\rm M}_2/{\rm M}_1= 10000:1$ ."89 These high mass ratio mergers certainly generate a different eravitational wave signal: they may even deserve a different source classification to separate it from the classical equal ass merger or extreme mass ratioral., These high mass ratio mergers certainly generate a different gravitational wave signal; they may even deserve a different source classification to separate it from the classical equal mass merger or extreme mass ratio.90. We use the rates found in our volune to extrapolate this eravitational wave signal over the IIubble Volume., We use the rates found in our volume to extrapolate this gravitational wave signal over the Hubble Volume.91 Naturally. this is highly seusitive to cosmic variance. and ten runs selected from an initial huger volume are planned to mitigate cosmic variance i our merecr rate estimates.," Naturally, this is highly sensitive to cosmic variance, and ten runs selected from an initial larger volume are planned to mitigate cosmic variance in our merger rate estimates."92 Towever. these s11all vole cosmological simulations are extremely nonlinear. auc as such. are computationally expensive.," However, these small volume cosmological simulations are extremely nonlinear, and as such, are computationally expensive."93" We are publishing the preliminary results frou, our first. smaller scale. siuulation i order to draw the attention of the conunuunitv to the interesting possibility. that the asseniblv of the lowest mass SMDBIIS may involve these lieh mass ratio mergers aud be a significant contributor to the LISA detectable event rate."," We are publishing the preliminary results from our first, smaller scale, simulation in order to draw the attention of the community to the interesting possibility that the assembly of the lowest mass SMBHs may involve these high mass ratio mergers and be a significant contributor to the LISA detectable event rate."94 We review the details of our simulation iu section 2 and in section 3. we describe how to caleulate the eravitational wave signal from two mereie black holes.," We review the details of our simulation in section 2 and in section 3, we describe how to calculate the gravitational wave signal from two merging black holes."95 We discuss our results aud implications in section [., We discuss our results and implications in section 4.96 In 7. we performed a ligh-+vesolutiou “Zzoom-i cosmological N-body simulation of a comoving section of a ACDAM universe (Q\;=0.3. O4 20.7. 040.9 and hi-0.7) from z=19 to z=0.," In \citet{KHB:2008bhgrowth}, , we performed a high-resolution 'zoom-in' cosmological N-body simulation of a comoving section of a $\Lambda$ CDM universe $\Omega_{\rm M}$ =0.3, $\Omega_{\Lambda}$ =0.7, $\sigma_8$ =0.9 and h=0.7) from $=49$ to $=0$."97 The high resolution region was a box 10 5h! Mpe on each side. for a total comoving volume of 1000 5.7 Mpc?.," The high resolution region was a box 10 $h^{-1}$ Mpc on each side, for a total comoving volume of 1000 $h^{-3}$ $^3$."98 The vole is designed to provide oue possible evolutionary path for a region like our Local Croup. aud as such. it should contain supermassive black holes ou the light end of the supermassive black hole uass Spectruni.," The volume is designed to provide one possible evolutionary path for a region like our Local Group, and as such, it should contain supermassive black holes on the light end of the supermassive black hole mass spectrum."99 Our mass resolution is 8.85«10M... and our spatial resolution is 2 kpc.," Our mass resolution is $8.85 \times 10^5 \, {\rm M}_\odot$, and our spatial resolution is 2 kpc."100 After the simulation is complete. we identified alos with at least 32 particles using P-Croupfiuder (2) and seed black holes those halos in the appropriate mass and redshift range to lost Pop III stars.," After the simulation is complete, we identified halos with at least 32 particles using P-Groupfinder \citep{pgroupfinder} and seed black holes those halos in the appropriate mass and redshift range to host Pop III stars."101 Note that we are using WALAPS (7) cosinological parameters in this study to compare with our previous work: however. at the time of this papers submission. zoonriu sinulations of several small volumes are underway with WALAPS parameters to better explore cosmic variance.," Note that we are using WMAP3 \citep{spergel:07} cosmological parameters in this study to compare with our previous work; however, at the time of this paper's submission, 'zoom-in' simulations of several small volumes are underway with WMAP5 parameters to better explore cosmic variance."102 This will allow us to better pin down the rate of these uew eravitational wave SOULCOS., This will allow us to better pin down the rate of these new gravitational wave sources.103 Iu our hvbrid method. we combine the dark matter halo merger trees obtained m muuerical simulations witli an analytical treatinent of the plhisical processes that arise m the dvnamics of galaxy aud black hole merecrs.," In our hybrid method, we combine the dark matter halo merger trees obtained in numerical simulations with an analytical treatment of the physical processes that arise in the dynamics of galaxy and black hole mergers."104 Since some of the processes are ill-constrained. we probe the effect of different black hole erowth recipes on the final black hole mass function.," Since some of the processes are ill-constrained, we probe the effect of different black hole growth recipes on the final black hole mass function."105 Du general. we asune cach dark matter halo is described bv au NEW profile (?).. aud include the effects of dviauical friction aud merecr- eas accretion. outo the SMDITL. as wellas the SMIDIT σος itself.," In general, we assume each dark matter halo is described by an NFW profile \citep{nfw:97}, and include the effects of dynamical friction and merger-induced gas accretion onto the SMBH, as wellas the SMBH merger itself."106 This hiybrid approach generates. for cach recipe. a census of the number aud mass ratio of the niergers in our volume at cach redshift.," This hybrid approach generates, for each recipe, a census of the number and mass ratio of the mergers in our volume at each redshift."107General relativity. a cornerstone of physics. is arguably one of our greatest intellectual achievements.,"General relativity, a cornerstone of physics, is arguably one of our greatest intellectual achievements."108 Lt is not only elegant ancl physically motivated. but it makes a whole host of predictions including gravitational waves. the anomalous precession of Mercury and the dellection of light.all of which have been verified.," It is not only elegant and physically motivated, but it makes a whole host of predictions including gravitational waves, the anomalous precession of Mercury and the deflection of light–all of which have been verified."109 Further still. clue to its weak vet accumulative nature. gravitation is a force of the large scale.," Further still, due to its weak yet accumulative nature, gravitation is a force of the large scale."110 And so cosmology. being the study of the Universe. is a field performed. within the formalism of this theory.," And so cosmology, being the study of the Universe, is a field performed within the formalism of this theory."111 Lt is in this way that Cosmology, It is in this way that Cosmology112"is constrained primarily by the πας data, with a small contribution from the Integrated Sachs-Wolfe (ISW) effect in the CMB.","is constrained primarily by the $f_{\rm gas}$ data, with a small contribution from the Integrated Sachs-Wolfe (ISW) effect in the CMB."113 A simple exercise provides further insight into how the CMB data help in constraining dark energy., A simple exercise provides further insight into how the CMB data help in constraining dark energy.114" For this, we re-examine the constraints in the wo—wa plane obtained from the fgas+CMB data; the 68 and 95 per cent confidence contours for the default model with 5 per cent allowances are shown (dashed curves) in Figure 6.."," For this, we re-examine the constraints in the $w_{\rm 0}-w_{\rm a}$ plane obtained from the $f_{\rm gas}$ +CMB data; the 68 and 95 per cent confidence contours for the default model with 5 per cent allowances are shown (dashed curves) in Figure \ref{fig:cmbede}. ."115" The combination of feas+CMB data provides tight constraints on QyA?, Qanh? and J, (driven primarily by the CMB data), and on A (driven by the combination of both data sets)."," The combination of $f_{\rm gas}$ +CMB data provides tight constraints on $\Omega_{\rm b}h^2$, $\Omega_{\rm116 dm}h^2$ and $l_{\rm a}$ (driven primarily by the CMB data), and on $h$ (driven by the combination of both data sets)."117" Using these constraints as priors, we examine the constraints in the Wo—Wa plane that can be obtained from the fas data alone; the results are shown as the red, solid curves in Figure 6.."," Using these constraints as priors, we examine the constraints in the $w_{\rm118 0}-w_{\rm a}$ plane that can be obtained from the $f_{\rm gas}$ data alone; the results are shown as the red, solid curves in Figure \ref{fig:cmbede}."119" We see that the priors encompass some of the CMB constraining power, in particular in defining the characteristic upper boundary in the two—wa plane."," We see that the priors encompass some of the CMB constraining power, in particular in defining the characteristic upper boundary in the $w_{\rm 0}-w_{\rm a}$ plane."120" However, they do not contain the full information on, e.g., the covariance of Qy?, Qa,h? and h which is also important in constraining dark energy at later times."," However, they do not contain the full information on, e.g., the covariance of $\Omega_{\rm b}h^2$, $\Omega_{\rm dm}h^2$ and $h$ which is also important in constraining dark energy at later times."121 We note that the prior on /4 provides a tight constraint on the curvature., We note that the prior on $l_{\rm a}$ provides a tight constraint on the curvature.122" The blue dotted curves in Figure 6 show the constraints obtained from the fas data alone, using only the priors on Q,h? and h and assuming flatness."," The blue dotted curves in Figure \ref{fig:cmbede} show the constraints obtained from the $f_{\rm gas}$ data alone, using only the priors on $\Omega_{\rm b}h^2$ and $h$ and assuming flatness."123 The XSZ technique provides a complementary and independent experiment to measure dark energy., The XSZ technique provides a complementary and independent experiment to measure dark energy.124" Although the inclusion of constraints from the XSZ experiment leads to only modest formal improvements in the FoM with respect to the results for the fz4,4-CMB data (Table 4;; as can be expected given the relatively weak dependence on dark energy in equation 21), it is important to note that the XSZ experiment relies on different assumptions and has different systematic uncertainties."," Although the inclusion of constraints from the XSZ experiment leads to only modest formal improvements in the FoM with respect to the results for the $f_{\rm gas}+$ CMB data (Table \ref{tab:models}; as can be expected given the relatively weak dependence on dark energy in equation 21), it is important to note that the XSZ experiment relies on different assumptions and has different systematic uncertainties."125" In particular, the XSZ experiment is independent of assumptions regarding hydrostatic equilibrium, the depletion factor, and the stellar mass fraction."," In particular, the XSZ experiment is independent of assumptions regarding hydrostatic equilibrium, the depletion factor, and the stellar mass fraction."126" Thus, the combination of data from the fas and XSZ techniques can help to ensure robustness in the results."," Thus, the combination of data from the $f_{\rm gas}$ and XSZ techniques can help to ensure robustness in the results."127" In principle, the inclusion of XSZ data can also allow some of the priors in the [κας experiment to be relaxed."," In principle, the inclusion of XSZ data can also allow some of the priors in the $f_{\rm gas}$ experiment to be relaxed."128" We have examined the ability of a future X-ray observatory, with capabilities similar to those planned for Constellation-X, to constrain dark energy via the [κας experiment."," We have examined the ability of a future X-ray observatory, with capabilities similar to those planned for Constellation-X, to constrain dark energy via the $f_{\rm gas}$ experiment."129" We find that feas measurements for a sample of 500 hot (kT25002 5keV), X-ray bright, dynamically relaxed clusters, with a precision of ~5 per cent, can be used to constrain dark energy with a FoM of 15—40."," We find that $f_{\rm gas}$ measurements for a sample of 500 hot $kT_{2500}\gsim5$ keV), X-ray bright, dynamically relaxed clusters, with a precision of $\sim 5$ per cent, can be used to constrain dark energy with a FoM of $15-40$."130" These constraints are comparable to those predicted by the DETF for other leading, planned (DETF Stage IV) dark energy experiments."," These constraints are comparable to those predicted by the DETF for other leading, planned (DETF Stage IV) dark energy experiments."131" We also find that, for the fea; experiment, the FoM can be boosted up by at least ~40 per cent by selecting an optimal redshift distribution of suitable clusters on which to carry out the fgas observations."," We also find that, for the $f_{\rm gas}$ experiment, the FoM can be boosted up by at least $\sim 40$ per cent by selecting an optimal redshift distribution of suitable clusters on which to carry out the $f_{\rm gas}$ observations."132" Interestingly, the optimal redshift distribution of feas measurments appears to be shifted towards low redshifts."," Interestingly, the optimal redshift distribution of $f_{\rm gas}$ measurments appears to be shifted towards low redshifts."133" As discussed in the text, a future [κας experiment will need to be preceded by a large X-ray or SZ clustersurvey that will find hot, X-ray luminous clusters out to high redshifts."," As discussed in the text, a future $f_{\rm gas}$ experiment will need to be preceded by a large X-ray or SZ clustersurvey that will find hot, X-ray luminous clusters out to high redshifts."134 A survey such as that planned with the Spectrum- mission should find several thousand of such clusters., A survey such as that planned with the Spectrum-RG/eROSITA mission should find several thousand of such clusters.135 Short ‘snapshot’ follow-up observations of the, Short `snapshot' follow-up observations of the136In order to constrain the dvnamical model. it is crucial to have a good estimate of the radius of the companion stu.,"In order to constrain the dynamical model, it is crucial to have a good estimate of the radius of the companion star."137 However. customary methods of determining (this radius fail because (he Cvenus N-1 svstem does not exhibit eclipses nor does the companion star fill its Roche equipotential lobe.," However, customary methods of determining this radius fail because the Cygnus X-1 system does not exhibit eclipses nor does the companion star fill its Roche equipotential lobe."138 We obtain the required estimate of the stellar radius as we have done previously in our study of LMC. X-1 (Oroszetal.2009)., We obtain the required estimate of the stellar radius as we have done previously in our study of LMC X-1 \citep{oro+2009}.139. The radius. which critically depends on distance. additionally. depends on the apparent magnitude of the O-(vpe star and interstellar extinction. ancl also on the effective stellar temperature ancl corresponding bolometrie correction.," The radius, which critically depends on distance, additionally depends on the apparent magnitude of the O-type star and interstellar extinction, and also on the effective stellar temperature and corresponding bolometric correction."140" The absolute magnitude of (he star is Ma,=A 0.11.1,. where A is the apparent A-band magnitude. BC, is the bolometric correction lor the A-band. D is the distance. and Ay is the extinction in the V-baud."," The absolute magnitude of the star is $M_{\rm141abs}=K+BC_K(T_{\rm eff},g)-(5\log D-5) - 0.11A_V$ , where $K$ is the apparent $K$ -band magnitude, $BC_K$ is the bolometric correction for the $K$ -band, $D$ is the distance, and $A_V$ is the extinction in the $V$ -band."142 The luminosity and radius of the star in solar units are L=10MMos15D and Tay). respectively.," The luminosity and radius of the star in solar units are $L=10^{-0.4(M_{\rm abs}-4.71)}$ and $R=\sqrt{L(5770/T_{\rm eff})^4}$ , respectively."143 In computing these quantities. we use D=1.86um kpc (Reid and à A-band apparent magnitude of ἐν=6.502:0.02 (Skrutskieetal.2006).. which minimizes the effects of interstellar extinction.," In computing these quantities, we use $D = 1.86_{-0.11}^{+0.12}$ kpc \citep{rei+2010} and a $K$ -band apparent magnitude of $K=6.50\pm 0.02$ \citep{skr+2006}, which minimizes the effects of interstellar extinction."144 For the A-band extinction. we adopt LE(B—V)-lllz.03and Ry=3.0240.03 (e.g.Ay=3.35.Caballero-Nievesetal.2009) and use the standard extinction law (Cardelliοἱal.1939).," For the $K$ -band extinction, we adopt $E(B-V)=1.11\pm 0.03$ and $R_V=3.02\pm 0.03$ \citep[e.g.\ $A_V=3.35$,][]{cab+2009} and use the standard extinction law \citep{car+1989}."145. The bolometrie corrections for the {ρα were computed using the OSTAR2002 grid of models with solar metallicity 2003)., The bolometric corrections for the $K$ -band were computed using the OSTAR2002 grid of models with solar metallicity \citep{lan+2003}.146. We note that the A-band bolometric corrections for the solar metallicity models ancl (he models for hall-solar metallicity differ only by 0.02 dex LLanz private communication). so our results are nol sensitive (o the metallicity.," We note that the $K$ -band bolometric corrections for the solar metallicity models and the models for half-solar metallicity differ only by 0.02 dex Lanz private communication), so our results are not sensitive to the metallicity."147 Figure 1. shows the derived radius and luminosity of the star as à function of its assumec temperature in the range28.000<Tig34.000 Ix. For Tog=28.000. the radius is 1...," Figure \ref{plotraddist} shows the derived radius and luminosity of the star as a function of its assumed temperature in the range$28,000148\le T_{\rm eff}\le 34,000$ K. For $T_{\rm eff}=28,000$, the radius is $R_{\rm dist}=19.26\pm1490.98\,R_{\odot}$ ."150 As the temperature increases. the radius decreases rapidly at first. and then it plateaus midwav through the range. attaining a value of Raw=16.3440.8411. al Teg=34.000 Ix. Meanwhile. (he luminosity increases with temperature. rising [rom L—2.x10L. at Taj=28.000 K to L—3.2xLOL. at Tag=34.000 K. The elfective temperature of the companion star can be determined [rom a detailed analvsis of UV and optical line spectra (lerveroetal.1995:IxaritskavaNievesοἱal. 2009).," As the temperature increases, the radius decreases rapidly at first, and then it plateaus midway through the range, attaining a value of $R_{\rm dist}=16.34\pm1510.84\,R_{\odot}$ at $T_{\rm eff}=34,000$ K. Meanwhile, the luminosity increases with temperature, rising from $L=2.1\times 10^5\,L_{\odot}$ at $T_{\rm eff}=28,000$ K to $L=3.2\times 10^5\,L_{\odot}$ at $T_{\rm152eff}=34,000$ K. The effective temperature of the companion star can be determined from a detailed analysis of UV and optical line spectra \citep{her+1995,kar+2005,cab+2009}."153. ILowever. it is often diffieult to determine a precise temperature for O-(vpe stars owing to a correlation between (he effective temperature ζω ancl (he surface eravity parameter log g.," However, it is often difficult to determine a precise temperature for O-type stars owing to a correlation between the effective temperature $T_{\rm eff}$ and the surface gravity parameter $\log g$ ."154 Model aüimosphereswith slightly smaller values of Ziq ancl loggy eive spectra (that are very similar to those obtained for slishtly higher values of (hese parameters., Model atmosphereswith slightly smaller values of $T_{\rm eff}$ and $\log g$ give spectra that are very similar to those obtained for slightly higher values of these parameters.155"silicate, can be modified by addition of corundum dust.","silicate, can be modified by addition of corundum dust."156" The AK ARI, IRAS and MSX flux data at 12, 9, 14.65 and 8.28 im would exhibit the presence of corundum dust efficiently."," The $AKARI$ $IRAS$ and $MSX$ flux data at 12, 9, 14.65 and 8.28 $\mu$ m would exhibit the presence of corundum dust efficiently."157" Many 2CDs displayed in Figs 4, 5, 6 and 7 show that a mixture of silicate (80 by mass) and corundum (20 %)) can make a better fit with observations of O-rich stars rather than pure silicate."," Many 2CDs displayed in Figs 4, 5, 6 and 7 show that a mixture of silicate (80 by mass) and corundum (20 ) can make a better fit with observations of O-rich stars rather than pure silicate."158" When we change the abundance of corundum from 0 to 20 gradually, we find that the series of model tracks can cover a wide area of the most populated observation points fairly well for many 2CDs."," When we change the abundance of corundum from 0 to 20 gradually, we find that the series of model tracks can cover a wide area of the most populated observation points fairly well for many 2CDs."159 This effect is more clear on the 2CD using [12]—[25] versus K —[12] (the lower panel of Fig., This effect is more clear on the 2CD using $-$ [25] versus $K-$ [12] (the lower panel of Fig.160 6)., 6).161 Other 2CDs using the fluxes at 12 and 9 um also show the effect well., Other 2CDs using the fluxes at 12 and 9 $\mu$ m also show the effect well.162" The effect of corundum dust would be prominent for the 2CDs using the fluxes at 12, 9, 14.65 and 8.28 µπι. However, the effect of corundum abundance may not improve the fitting with observations for some 2CDs (e.g., the 2CD in the lower panel of Fig."," The effect of corundum dust would be prominent for the 2CDs using the fluxes at 12, 9, 14.65 and 8.28 $\mu$ m. However, the effect of corundum abundance may not improve the fitting with observations for some 2CDs (e.g., the 2CD in the lower panel of Fig."163 7)., 7).164" For C-rich stars, we use the optical constants of amorphous carbon (AMO) grains derived by Suh (2000) and the optical constants of a SiC grains by Péggourié (1988)."," For C-rich stars, we use the optical constants of amorphous carbon (AMC) grains derived by Suh (2000) and the optical constants of $\alpha$ SiC grains by Péggourié (1988)."165 The radii of spherical dust grains have been assumed to be 0.1 um uniformly., The radii of spherical dust grains have been assumed to be 0.1 $\mu$ m uniformly.166" We choose 10 µπι as the fiducial wavelength that sets the scale of the optical depth(rio) and perform the model calculations for seven optical depths (r19 = 0.01, 0.1, 1, 2,3, 5 and 7)."," We choose 10 $\mu$ m as the fiducial wavelength that sets the scale of the optical $\tau_{10}$ ) and perform the model calculations for seven optical depths $\tau_{10}$ $=$ 0.01, 0.1, 1, 2,3, 5 and 7)."167" For the central star, we assume that the"," For the central star, we assume that the"168We have worked out simple models of time dependent wind. devised to reproduce the eross features of the wind history predicted bv binary svstem evolutionary models.,"We have worked out simple models of time dependent wind, devised to reproduce the gross features of the wind history predicted by binary system evolutionary models."169" In all our wind models the rate αἱ which mass is lost by the binary svstem as a wind. AM... has a linear time dependence: M,=a— bl. with a. b (wo positive parameters. and the wind velocity. ey. remains constant with time."," In all our wind models the rate at which mass is lost by the binary system as a wind, $\dot{M}_{\rm w}$, has a linear time dependence: $\dot{M}_{\rm w} = a - bt$ , with $a$, $b$ two positive parameters, and the wind velocity, $v_{\rm w}$, remains constant with time."170 The wind can either be active until the moment of supernova explosion. as in models C (fast wind) and D (slow wind). or it can decrease to zero. leaving (he svstem in a phase in which there is no wind ejection until the explosion ensues. as in models A (fast wind) and D (slow wind).," The wind can either be active until the moment of supernova explosion, as in models C (fast wind) and D (slow wind), or it can decrease to zero, leaving the system in a phase in which there is no wind ejection until the explosion ensues, as in models A (fast wind) and B (slow wind)."171 The parameters for the models are given in table 1. and figure I.., The parameters for the models are given in table \ref{tbl-1} and figure \ref{fig-1}.172 We have studied the hvdrodvnanmical processes induced by the wind ejection: first. its interaction with the AM and. second. the formation ancl evolution of the SNR.," We have studied the hydrodynamical processes induced by the wind ejection: first, its interaction with the AM and, second, the formation and evolution of the SNR."173 The simulations were perlormed with a standard one dimensional hydrodyvnanmical code similar to (that described in Truelove&MclIxee(1999)... with the equations modified to include a source of mass and momentum at the center. to simulate wind ejection.," The simulations were performed with a standard one dimensional hydrodynamical code similar to that described in \citet{tm99}, with the equations modified to include a source of mass and momentum at the center, to simulate wind ejection."174 The ISM density was set to LO>7 @  3:in all cases., The ISM density was set to $10^{-24}$ g $^{-3}$ in all cases.175 The resulting AM density. profiles at the time of SN explosion are shown in ligure 2.., The resulting AM density profiles at the time of SN explosion are shown in figure \ref{fig-3}.176 The interaction of the wind ejected by the progenitor svstem aud (he surrounding ISM follows a mechanism very similar to that at work in SNI: high speed ejecta flow into a uniform. stationarv mediun and push it away resulting in a doublv-shocked structure with a contact discontinuitv between the wind ancl the ISM.," The interaction of the wind ejected by the progenitor system and the surrounding ISM follows a mechanism very similar to that at work in SNR: high speed ejecta flow into a uniform, stationary medium and push it away resulting in a doubly-shocked structure with a contact discontinuity between the wind and the ISM."177 A forward shock (prominent in all the profiles shown in figure 2)) propagates into the ISM. heating. compressing and accelerating it. while a reverse shock propagates inward. heating. compressing aud decelerating the inner windmaterial.," A forward shock (prominent in all the profiles shown in figure \ref{fig-3}) ) propagates into the ISM, heating, compressing and accelerating it, while a reverse shock propagates inward, heating, compressing and decelerating the inner windmaterial."178 The kinetic energies involved. however. are four to," The kinetic energies involved, however, are four to"179scheduling.,scheduling.180 Thev have contributed 3 new millisecond pulsars and 1G slow ones to the total number of the drif(-sean survey discoveries quoted above., They have contributed 3 new millisecond pulsars and 16 slow ones to the total number of the drift-scan survey discoveries quoted above.181 The millisecond pulsars P5R J1640--2224 and PSR JI713--074T have been svstematically monitored at Arecibo since the time of their discovery., The millisecond pulsars PSR J1640+2224 and PSR J1713+0747 have been systematically monitored at Arecibo since the time of their discovery.182 The Uiming models for these objects have been published by Wolszezan et al. (, The timing models for these objects have been published by Wolszczan et al. (1832000) and Camilo. Foster and Wolszezan (1994). respectively.,"2000) and Camilo, Foster and Wolszczan (1994), respectively."184 Our continuing observations discussed in this paper have established an accurate üming model for the third millisecond pulsar. PSR J110942313. positively verified (wo of the four unconfirmed slow pulsar candidates included in Foster et al. (," Our continuing observations discussed in this paper have established an accurate timing model for the third millisecond pulsar, PSR J1709+2313, positively verified two of the four unconfirmed slow pulsar candidates included in Foster et al. ("1851995). PSR J2151+2315 and PSR 21554-2813. and discovered two acditional slow pulsars. PSI. 18132-1322 and PSR J1908+2351.,"1995), PSR J2151+2315 and PSR 2155+2813, and discovered two additional slow pulsars, PSR 1813+1822 and PSR J1908+2351."186 Two of our seeminelv original discoveries. PSR J15494-21 and PSR J19064-16. have proven to be pulsars detected by other searches (Lorimer. private communication: Camilo 1995).," Two of our seemingly original discoveries, PSR J1549+21 and PSR J1906+16, have proven to be pulsars detected by other searches (Lorimer, private communication; Camilo 1995)."187 Finally. the slow pulsar candidates. PSR D18402-13 and PSR D22084-18. and the millisecond pulsar candidate. PSR. D17354-13. have not been confirmed.," Finally, the slow pulsar candidates, PSR B1840+13 and PSR B2208+18, and the millisecond pulsar candidate, PSR B1735+13, have not been confirmed."188 In (his paper. we present refined timing models for the millisecond pulsar PSR J11094-23. the pulsar PSI. J05384-2817 in the supernova remnant S147. and the sixteen confirmed slow pulsars discovered by (the Penn Stale/NRL surveys since 1991.," In this paper, we present refined timing models for the millisecond pulsar PSR J1709+23, the pulsar PSR J0538+2817 in the supernova remnant S147, and the sixteen confirmed slow pulsars discovered by the Penn State/NRL surveys since 1991."189 The model for PSR J1709+2313 includes a significant proper motion measurement., The model for PSR J1709+2313 includes a significant proper motion measurement.190 For all pulsars listed in Foster et al. (, For all pulsars listed in Foster et al. (1911995). both the spin parameters and the (imine positions have been significantly improved.,"1995), both the spin parameters and the timing positions have been significantly improved."192 Dining parameters for the (wo new slow pulsars are published here for the first lime., Timing parameters for the two new slow pulsars are published here for the first time.193 We also present the results of single pulse observations of the brightest objects in the sample., We also present the results of single pulse observations of the brightest objects in the sample.194" These include (he unusual ""bursting pulsar. PSR J175242-2359. two new drifting subpulse pulsars. PSR. J1649+2533 and PSR J2155+2813. and a new mode changing pulsar. PSR J17464-2540."," These include the unusual “bursting” pulsar, PSR J1752+2359, two new drifting subpulse pulsars, PSR J1649+2533 and PSR J2155+2813, and a new mode changing pulsar, PSR J1746+2540."195 Our observations and cata analysis are described in Section 2., Our observations and data analysis are described in Section 2.196 In Section 3. we give the details of the new timing models for the millisecond. pulsar PSR J17094-2313.. PSR JO538+2817. the pulsar located inside the supernova remnant 5147. and [or the 16 slow pulsars.," In Section 3, we give the details of the new timing models for the millisecond pulsar PSR J1709+2313, PSR J0538+2817, the pulsar located inside the supernova remnant S147, and for the 16 slow pulsars."197 Section 4 is devoted to the single pulse analvsis of three pulsars from this sample and our conclusions are given in Section 5., Section 4 is devoted to the single pulse analysis of three pulsars from this sample and our conclusions are given in Section 5.198 The pulse timing and the single pulse observations discussed in (his paper have been made with the 305-m Arecibo radiotelescope using the cdualcirenlar polarization receiving svstems al 430 MIIz and 1400 MIIz and the Penn State Pulsar Machine (PSPAI)., The pulse timing and the single pulse observations discussed in this paper have been made with the 305-m Arecibo radiotelescope using the dual-circular polarization receiving systems at 430 MHz and 1400 MHz and the Penn State Pulsar Machine (PSPM).199 Timing observations made at Arecibo before 1995 with the 40 MlIIz correlation spectrometer as a pulsar backend are discussed in Foster οἱ al. (, Timing observations made at Arecibo before 1995 with the 40 MHz correlation spectrometer as a pulsar backend are discussed in Foster et al. (2001995) ancl Cacwell (1997).,1995) and Cadwell (1997).201Variability was very carly established as a trade mark of Active Galactic Nuclei (AGN).,Variability was very early established as a trade mark of Active Galactic Nuclei (AGN).202 Variability has also been used. as à key tool to derive physical properties of AGN: characteristic time scales were used to infer sizes of the emitting regions. lags between the ionising continuum and the line response have been used to determine Black Hole masses (A). multi-wavelength light. curves have been used to study the ecometrical and physical connections between the different regions around the central engine.," Variability has also been used as a key tool to derive physical properties of AGN: characteristic time scales were used to infer sizes of the emitting regions, lags between the ionising continuum and the line response have been used to determine Black Hole masses $M$ ), multi-wavelength light curves have been used to study the geometrical and physical connections between the different regions around the central engine."203 The scenario accepted. until recently for the interplay of the emitting regions is that variability is driven by the emission from the X-rav corona located close to the central Black Hole (c.g. Collin-Soulfrin. 1991: Ixrolik et 11991: Clavel et al.," The scenario accepted until recently for the interplay of the emitting regions is that variability is driven by the emission from the X-ray corona located close to the central Black Hole (e.g., Collin-Souffrin, 1991; Krolik et 1991; Clavel et al.,"204 1992: Collier et 11999: Cackett et 22007)., 1992; Collier et 1999; Cackett et 2007).205 The negligible optical inter-band lags was carly evidence that pointed. towards reprocessing of high-energy photons bv the accretion disc. where the characteristic distances between the cilferent emitting regions correspond to the ight travelling time (Collin-Soullrin.. 1991: WKerolik ct 11991: Clavel et aL.," The negligible optical inter-band lags was early evidence that pointed towards reprocessing of high-energy photons by the accretion disc, where the characteristic distances between the different emitting regions correspond to the light travelling time (Collin-Souffrin, 1991; Krolik et 1991; Clavel et al.,"206 1992)., 1992).207 Later on. the measurement of short lags between the X-ray. emission and the optical. anc he leading of the X-ravs whenever significant lags were determined (es. IEdelson et al..," Later on, the measurement of short lags between the X-ray emission and the optical, and the leading of the X-rays whenever significant lags were determined (e.g., Edelson et al.,"208. 1996: Wanders et al.," 1996; Wanders et al.,"209 1997: Shemumer et al..," 1997; Shemmer et al.,"210 2001: Desroches ct al..," 2001; Desroches et al.,"211 2006). lent support o this picture (see also Nandra et al..," 2006), lent support to this picture (see also Nandra et al.,"212. 2000 for evidence supporting reprocessing from the correlation between the UV flux and the spectral shape of the X-ray emission). since or intrinsic cise variability shorter wavelengths should. lag onger wavelength. emission by long (viscous) time-scales.," 2000 for evidence supporting reprocessing from the correlation between the UV flux and the spectral shape of the X-ray emission), since for intrinsic disc variability shorter wavelengths should lag longer wavelength emission by long (viscous) time-scales."213 Also. light curves showed that the amplitude of the X-ray variations was much larger than that seen in the optical. which can be explained by the damping of the signal during the disc reprocessing.," Also, light curves showed that the amplitude of the X-ray variations was much larger than that seen in the optical, which can be explained by the damping of the signal during the disc reprocessing."214 However. a full picture of the interplay between the X-ray corona and the disc might. not be complete just vet. as new evidence seems to suggest that long termi optical variability is. driven by accretion," However, a full picture of the interplay between the X-ray corona and the disc might not be complete just yet, as new evidence seems to suggest that long term optical variability is driven by accretion"215The interesting transition hinting for an instability was reported by Homanetal.(2001).,The interesting transition hinting for an instability was reported by \cite{homan01}.216" In the MJD 51,254 observation, when the source was still very bright, the luminosity suddenly increased without a change in the color."," In the MJD 51,254 observation, when the source was still very bright, the luminosity suddenly increased without a change in the color."217" Whether indeed this single transition hints for the radiation pressure instability or not, the source likley defines the lower limit for the radiation pressure instability to operate."," Whether indeed this single transition hints for the radiation pressure instability or not, the source likley defines the lower limit for the radiation pressure instability to operate."218 In Table 1 we do not see sources which have very large Eddington ratio and are stable against the radiation pressure instability., In Table \ref{tab:binaries} we do not see sources which have very large Eddington ratio and are stable against the radiation pressure instability.219" As we mentioned above, GRS 1915--105 is a good example of showing outbursts even at Eddington ratio close to 1 so it seems we have no upper limit for the Eddington ratio in the case of radiation pressure instability."," As we mentioned above, GRS 1915+105 is a good example of showing outbursts even at Eddington ratio close to 1 so it seems we have no upper limit for the Eddington ratio in the case of radiation pressure instability."220" Thus, observationally, the radiation pressure instability should operate between the Eddington ratio 0.15 up to 1 or more."," Thus, observationally, the radiation pressure instability should operate between the Eddington ratio 0.15 up to 1 or more."221 Comparing this with the several theoretical possibilities plotted in Fig., Comparing this with the several theoretical possibilities plotted in Fig.222 1 we can draw certain conclusions., \ref{fig:topo} we can draw certain conclusions.223" First, only the viscosity prescription a\/PeasProt is consistent with the lower limit for the radiation pressure instability, as the unstable region then extends from the Eddington 0.16 up."," First, only the viscosity prescription $\alpha \sqrt{P_{\rm gas}P_{\rm tot}}$ is consistent with the lower limit for the radiation pressure instability, as the unstable region then extends from the Eddington 0.16 up."224 The prescription aPiot would allow instability to operate at too low luminosity., The prescription $\alpha P_{\rm tot}$ would allow instability to operate at too low luminosity.225" Second, too efficient cooling by the jet is also ruled out."," Second, too efficient cooling by the jet is also ruled out."226 The cooling operates similarly in both cases of viscosity parameterization and stabilize the disc., The cooling operates similarly in both cases of viscosity parameterization and stabilize the disc.227" For the adopted values of the jet efficiency parameter, the disc is stable for Eddington ratio above 0.22, which is clrearly inconsistent with observations."," For the adopted values of the jet efficiency parameter, the disc is stable for Eddington ratio above 0.22, which is clrearly inconsistent with observations."228" Therefore, the parameter A is Eq.1 of the disk-jet coupling must be significantly lower than this exemplary value of A—25."," Therefore, the parameter $A$ is \ref{eq:jet} of the disk-jet coupling must be significantly lower than this exemplary value of $A=25$."229" However the jet is by no means excluded and still can carry a substantial energy, because in the case of equipartition between the disk and jet radiation, for the Eddington accretion rate m=1 the jet coupling constant equal to A=1 would be enough."," However the jet is by no means excluded and still can carry a substantial energy, because in the case of equipartition between the disk and jet radiation, for the Eddington accretion rate $\dot m =1$ the jet coupling constant equal to $A=1$ would be enough."230" The parameterization a,/PsasPio has an additional advantage of reducing the outburst amplitude in comparison to aPiot.", The parameterization $\alpha \sqrt{P_{\rm gas}P_{\rm tot}}$ has an additional advantage of reducing the outburst amplitude in comparison to $\alpha P_{\rm tot}$.231" Most of the candidate sources for radiation pressure instability show rather low to moderate amplitudes, from factor 2 to 20."," Most of the candidate sources for radiation pressure instability show rather low to moderate amplitudes, from factor 2 to 20."232" Only one source - GS 2023+338 - shows huge outbursts, with the factor of 500 brightenings in timescales of 60 seconds."," Only one source - GS 2023+338 - shows huge outbursts, with the factor of 500 brightenings in timescales of 60 seconds."233 in’tZandetal.(1992) interpreted this short timescale variability as caused by variable absorption., \cite{zand92} interpreted this short timescale variability as caused by variable absorption.234 The behaviour of this source is exceptional and puzzling., The behaviour of this source is exceptional and puzzling.235" When studying the instabilities in the supermassive black hole environment, we usually cannot directly observe a duty cycle of a one single object, since the black hole masses are large and the expected timescales are very long."," When studying the instabilities in the supermassive black hole environment, we usually cannot directly observe a duty cycle of a one single object, since the black hole masses are large and the expected timescales are very long."236" Instead, the statistical studies are useful here and we can find an evidence for the source episodic activity (e.g. Czernyetal. (2009)))."," Instead, the statistical studies are useful here and we can find an evidence for the source episodic activity (e.g. \cite{czerny09}) )."237" However, the exceptional object is NGC 4395, with the black hole mass of 3.6x10° (Petersonetal. (2005)))."," However, the exceptional object is NGC 4395, with the black hole mass of $3.6 \times 10^{5} M_{\odot}$ \cite{peterson05}) )."238" In this source, in principle, we Mocould observe the variability due to radiation pressure instability."," In this source, in principle, we could observe the variability due to radiation pressure instability."239" As was shown by Czernyetal.(2009),, the outbursts for the central black hole of the mass 10""Mo, should last below 100 years, so for a mass 30 times smaller, the outbursts should last ~3 years!"," As was shown by \cite{czerny09}, the outbursts for the central black hole of the mass $10^{7} M_{\odot}$, should last below 100 years, so for a mass 30 times smaller, the outbursts should last $\sim 3$ years!"240 No such outbursts are observed., No such outbursts are observed.241" However, this fact is actually consistent with our expectations, since the Eddington ratio in this source is only 1.2x103."," However, this fact is actually consistent with our expectations, since the Eddington ratio in this source is only $1.2\times 10^{-3}$."242" The source is thus stable with both aP,o¢ and ay/PeasPtot mechanisms and provides no useful constraints for the parameterization of the viscous torque.", The source is thus stable with both $\alpha P_{\rm tot}$ and $\alpha \sqrt{P_{\rm gas}P_{\rm tot}}$ mechanisms and provides no useful constraints for the parameterization of the viscous torque.243 Significant constraints can be obtained from radio galaxies., Significant constraints can be obtained from radio galaxies.244" In case of the accretion discs in radio galaxies, the Eddington ratios can be estimated e.g. through the correlation with the broad line luminosities (2007)))."," In case of the accretion discs in radio galaxies, the Eddington ratios can be estimated e.g. through the correlation with the broad line luminosities \cite{dai07}) )."245" The FR I and FR II sources in this sample have low Eddington ratios, of 0.00975 and 0.0096, for FR I and FR II sources, respectively."," The FR I and FR II sources in this sample have low Eddington ratios, of 0.00975 and 0.0096, for FR I and FR II sources, respectively."246 Observations clearly show that these sources are stable against the radiation pressure instability since they form very large scale radio structures., Observations clearly show that these sources are stable against the radiation pressure instability since they form very large scale radio structures.247" In particular, the central engine of FR II galaxies must be operating in a continuous way for millions of years."," In particular, the central engine of FR II galaxies must be operating in a continuous way for millions of years."248 Fig., Fig.249 1 shows that their stability is consistent with theory if the heating is given by o4/Pza;Pi«., \ref{fig:topo} shows that their stability is consistent with theory if the heating is given by $\alpha \sqrt{P_{\rm gas}P_{\rm tot}}$.250" On the other hand, the FSRQ sources with compact radio structures tend to have larger Eddington ratios."," On the other hand, the FSRQ sources with compact radio structures tend to have larger Eddington ratios."251" These sources may in fact exhibit episodic activity and the small size of the structure is indicating a new episode, as proposed by Czernyetal. (2009)."," These sources may in fact exhibit episodic activity and the small size of the structure is indicating a new episode, as proposed by \cite{czerny09}."252". Therefore, it seems that the assumption of the ay/PeasPtot can accomodate the observational constraints both for Galactic sources and AGN."," Therefore, it seems that the assumption of the $\alpha \sqrt{P_{\rm gas}P_{\rm tot}}$ can accomodate the observational constraints both for Galactic sources and AGN."253 A typical value of the Eddington ratio found in the SDSS sample of quasars by Kellyetal.(2010) is 0.05 with a scatter of 0.4 dex., A typical value of the Eddington ratio found in the SDSS sample of quasars by \cite{kelly10} is 0.05 with a scatter of 0.4 dex.254 This is also large enough for the episodic activity caused by the radiation pressure instability., This is also large enough for the episodic activity caused by the radiation pressure instability.255" Possibly, the selection effect is in fact the reason why we detect only the sources in the active state: most of the sources in the quiescent state are too dim to be detectable."," Possibly, the selection effect is in fact the reason why we detect only the sources in the active state: most of the sources in the quiescent state are too dim to be detectable."256" There is also a possibility that active galaxies at high Eddington ratios, close to 1, are actually stable due to the stabilizing power of jet/outflow."," There is also a possibility that active galaxies at high Eddington ratios, close to 1, are actually stable due to the stabilizing power of jet/outflow."257 This mechanism seems not to work efficiently in Galactic sources but the relative jet power in accreting sources rises with the black hole mass:, This mechanism seems not to work efficiently in Galactic sources but the relative jet power in accreting sources rises with the black hole mass:258"of the work in this letter is for an observer located outside the initial jet opening. 0,My.","of the work in this letter is for an observer located outside the initial jet opening, $\theta_{\rm obs}>\theta_0$."259 We have considered three different Jet models of increasing sophistication: the simplest being a point source moving along the jet axis ($22.1). and the most sophisticated is 2D hydrodynamical simulation ($22.3).," We have considered three different jet models of increasing sophistication; the simplest being a point source moving along the jet axis 2.1), and the most sophisticated is 2D hydrodynamical simulation 2.3)."260" The basic qualitative features of the light-curves are similar in all three models. for 0,),.704."," The basic qualitative features of the light-curves are similar in all three models, for $\theta_{\rm obs}>\theta_0$."261 Moreover. the uniform jet model (model 2. 822.2) is in rough quantitative agreement with the hydro-model.," Moreover, the uniform jet model (model 2, 2.2) is in rough quantitative agreement with the hydro-model."262" We find that ""orphan"" optical afterglows associated with axis Jets can be observed up to a constant 0,5. rather than a constant 0,1,/( as suggested by Dalal et al. ("," We find that ""orphan"" optical afterglows associated with off-axis jets can be observed up to a constant $\theta_{\rm obs}$ rather than a constant $\theta_{\rm obs}/\theta_0$ as suggested by Dalal et al. ("2632002). if one assumes a constant energy in the jet. rather than a constant flux at the time of the jet break for an on-axis observer.,"2002), if one assumes a constant energy in the jet, rather than a constant flux at the time of the jet break for an on-axis observer."264 This implies that future surveys for orphan afterelows may provide valuable data for the the distribution of jet opening angles (; and the true event rate of GRBs., This implies that future surveys for orphan afterglows may provide valuable data for the the distribution of jet opening angles $\theta_0$ and the true event rate of GRBs.265 The orphan optical events discussed here can be identified from the initial rise during which the spectral slope is typically ο>0. followed by a decay. on a time scale of ~130 days. and may show a large degree of linear polarization (xz.1054).," The orphan optical events discussed here can be identified from the initial rise during which the spectral slope is typically $\beta > 0$, followed by a decay, on a time scale of $\sim 1-30$ days, and may show a large degree of linear polarization $\lesssim 40\%$ )."266 The detection of such orphan afterelows may provide a new line of evidence in favor of jetted outflows in GRBs., The detection of such orphan afterglows may provide a new line of evidence in favor of jetted outflows in GRBs.267 Recently Huang. Dai and Lu (2001) have considered another scenario (failed GRBs) for producing orphan afterglows: this would increase the detection rate of orphan afterglows.," Recently Huang, Dai and Lu (2001) have considered another scenario (failed GRBs) for producing orphan afterglows; this would increase the detection rate of orphan afterglows."268" A good monitoring ofoptical transients may help distinguish failed GRBs from jets seen at 0,>(y. and improve our understanding of them."," A good monitoring of optical transients may help distinguish failed GRBs from jets seen at $\theta_{\rm obs} > \theta_0$, and improve our understanding of them."269 We thank Mark Miller for generating the data for model 3, We thank Mark Miller for generating the data for model 3.270 This research was supported by grants NSF PHY 99-79985 and MCA93S025 (computational support). NSF grant PHY- (JG).Lyman Spitzer. Jr. Fellowship (AP). NAGS-8128 and MIT-292701 (SW).," This research was supported by grants NSF PHY 99-79985 and MCA93S025 (computational support), NSF grant PHY-0070928 (JG),Lyman Spitzer, Jr. Fellowship (AP), NAG5-8128 and MIT-292701 (SW)."271"acerete during this time. but at much reduced accretion rates of 210"" M. yr5 for the 3=0.001 calculation and 110 ""M. vr|! for the 3=0.005 calculation (measured between 20 and SO years after stellar core formation: Fig. 149.","accrete during this time, but at much reduced accretion rates of $2\times 10^{-5}$ $_\odot$ $^{-1}$ for the $\beta=0.001$ calculation and $1\times 10^{-5}$ $_\odot$ $^{-1}$ for the $\beta=0.005$ calculation (measured between 20 and 50 years after stellar core formation; Fig. \ref{convergence}) )."272 Although the rate of convergence with increasing resolution is relatively slow. strong outflows are launched in all cases (see Appendix A).," Although the rate of convergence with increasing resolution is relatively slow, strong outflows are launched in all cases (see Appendix A)."273 The slow convergence is due to the strong interplay between the energy released by the formation of the stellar core. and the launching of the outflow which decreases accretion and. thus. reduces the energy that is released.," The slow convergence is due to the strong interplay between the energy released by the formation of the stellar core, and the launching of the outflow which decreases accretion and, thus, reduces the energy that is released."274 If the resolution is poor. the stellar core accretes more material before the energy released feeds back and manages to stop the accretion. whereas with higher resolution. the accretion of a small amount of mass can feed thermal energy into the infalling material more quickly and. thus. inhibit further accretion.," If the resolution is poor, the stellar core accretes more material before the energy released feeds back and manages to stop the accretion, whereas with higher resolution, the accretion of a small amount of mass can feed thermal energy into the infalling material more quickly and, thus, inhibit further accretion."275 As mentioned in Section 3.2.. the introduction of radiative transfer and a realistic equation of state increases the lifetimes of the first core phase compared to that obtained using our barotropic equation of state by factors of 1.5—3 (Fig. 5).," As mentioned in Section \ref{rapidly_rotating}, the introduction of radiative transfer and a realistic equation of state increases the lifetimes of the first core phase compared to that obtained using our barotropic equation of state by factors of 1.5–3 (Fig. \ref{first_core_time}) )."276 This is important. because the longer the first core phase lasts. the more easy it will be to observe.," This is important, because the longer the first core phase lasts, the more easy it will be to observe."277 Thelifetimes obtained using radiative transfer range from es400 years (with no rotation) to =3000 years (for 3= 0.01). whereas using the barotropic equation of state they ranged from zz100 to z1500 years.," Thelifetimes obtained using radiative transfer range from $\approx 400$ years (with no rotation) to $\approx 3000$ years (for $\beta=0.01$ ), whereas using the barotropic equation of state they ranged from $\approx 100$ to $\approx 1500$ years."278 As discussed above. the difference is due to the higher temperatures Cand thus higher pressures) that are obtained using the realistic physics rather than the barotropic equation of state. which slows the evolution towards the second collapse phase.," As discussed above, the difference is due to the higher temperatures (and thus higher pressures) that are obtained using the realistic physics rather than the barotropic equation of state, which slows the evolution towards the second collapse phase."279 This lengthening of the lifetimes with radiation hydrodynamics compared to barotropic calculations was also seen by ?.., This lengthening of the lifetimes with radiation hydrodynamics compared to barotropic calculations was also seen by \cite{Tomidaetal2010a}. .280 The lengthening occurs for both the non-rotating first cores.," The lengthening occurs for both the non-rotating first cores,"281by the computed model light curve.,by the computed model light curve.282 However. the theoretical light curve fit becomes worse for later phases.," However, the theoretical light curve fit becomes worse for later phases."283 The luminosity of the model light curve seems to rise again at around day 25 after the explosion., The luminosity of the model light curve seems to rise again at around day 25 after the explosion.284 Up to day 45. the model light curve has à second bump. which ts not observed in the light curves of SN 1999ee and SN 2002bo.," Up to day 45, the model light curve has a second bump, which is not observed in the light curves of SN 1999ee and SN 2002bo."285 In the infrared I band. the decline after the maximum phase is missing. as shown in Fig. 16..," In the infrared I band, the decline after the maximum phase is missing, as shown in Fig. \ref{fig:lc_i}."286 As in the R band. the rise in the beginning and maximum are well represented in the model light curve.," As in the R band, the rise in the beginning and maximum are well represented in the model light curve."287 However. at maximum. the luminosity of the SN la model light curve rises further. which is not seen in the observed light curves of SN 2002bo and 1999ee.," However, at maximum, the luminosity of the SN Ia model light curve rises further, which is not seen in the observed light curves of SN 2002bo and 1999ee."288 Around day 30. the difference between model and observed light curve in the I band are about | mag.," Around day 30, the difference between model and observed light curve in the I band are about 1 mag."289 Up to day 50. the model light curve declines. while the observed light curves show their second maximum around 40 days after explosion.," Up to day 50, the model light curve declines, while the observed light curves show their second maximum around 40 days after explosion."290 In Fig. 17..," In Fig. \ref{fig:lc_u-b},"291 we compared the relation U-B of the theoretical and observed light curves., we compared the relation U-B of the theoretical and observed light curves.292 The same comparision for the B-V relation is shownin Fig. 18.., The same comparision for the B-V relation is shownin Fig. \ref{fig:lc_b-v}. .293Fig.,Fig.294 9 (top panel) the CO fractional abundance is ~1077. high enough to lead a significant destruction ofN4H.," \ref{modelfigs}~ (top panel) the CO fractional abundance is $\sim10^{-4}$, high enough to lead a significant destruction of."295. Contrary to the case of aandHCO the increase of the ffractional abundance is slow. which indicates that gas phase chemistry dominates over pure evaporation.," Contrary to the case of and, the increase of the fractional abundance is slow, which indicates that gas phase chemistry dominates over pure evaporation."296 This is simply a consequence of the fact that most of the ammonia is still locked in water ice: in fact. as shown by the experiments by Collings et al. (," This is simply a consequence of the fact that most of the ammonia is still locked in water ice; in fact, as shown by the experiments by Collings et al. ("2972004). ammonia is released back into the gas phase only if temperatures of ~ 100-120 K are reached (Viti et al.,"2004), ammonia is released back into the gas phase only if temperatures of $\sim$ 100-120 K are reached (Viti et al."298 2004)., 2004).299 The maximum value. of 4.11075 (or N(NH3))3.2x10? 2). is reached at /~1.3xIO? yr. and then it drops until it reaches à constant abundance of 1x10? (or a column density of ~8.2x10 em).," The maximum value, of $\times10^{-8}$ (or $\simeq3.2\times10^{15}$ ), is reached at $t\simeq1.3\times10^5$ yr, and then it drops until it reaches a constant abundance of $\times10^{-9}$ (or a column density of $\sim8.2\times10^{13}$ )."300 This behavior. tthe decrease in the ffractional abundance. takes place through reactions of the mmolecule with the ions C andHCO™.," This behavior, the decrease in the fractional abundance, takes place through reactions of the molecule with the ions $^+$ and."301. In this situation. the C fractional abundance increases at late stages. while the fractional abundance of iis More or less constant. which is consistent with the fact that iis produced from the destruction of ((see above) and destroyed through reactions withs.," In this situation, the $^+$ fractional abundance increases at late stages, while the fractional abundance of is more or less constant, which is consistent with the fact that is produced from the destruction of (see above) and destroyed through reactions with."302. Regarding the western and eastern cores. which were modeled assuming the same core size (~0.04 pc) and same temperature (the maximum temperature is 25 K). we found that the fractional abundance of increases moderately with time. while the fractional abundances of CO.N»..N2H7.. and rise considerably during the initial stages due to desorption effects. similar to the case of the the central core model.," Regarding the western and eastern cores, which were modeled assuming the same core size $\sim0.04$ pc) and same temperature (the maximum temperature is 25 K), we found that the fractional abundance of increases moderately with time, while the fractional abundances of CO, and rise considerably during the initial stages due to desorption effects, similar to the case of the the central core model."303 At rx4x10° yr. παπά hhave constant abundance. around ~2x107!° (N(N2H7))=4x107 emp and ~1.8x107? (N(HCO53.6x10 oor »z9x10!!cm. assuming 240) for aandHCO.. respectively. while the ffractional abundance increases. up to ~5.2x107% (N(NH3)=[x10% em)» at £=~13x10° yr and then it remains roughly constant too.," At $t\simeq4\times10^3$ yr, and have constant abundance, around $\sim2\times10^{-10}$ $\simeq4\times10^{12}$ ) and $\sim1.8\times10^{-9}$ $\simeq3.6\times10^{13}$ or $\times10^{11}$, assuming ]=40) for and, respectively, while the fractional abundance increases, up to $\sim5.2\times10^{-8}$ $\simeq1\times10^{15}$ ) at $t\simeq1.3\times10^{5}$ yr, and then it remains roughly constant too."304 For these cores. the values obtained from the chemical modeling are in agreement (within a factor of 2 in the case ofN>H™.. and a factor of 4 in the case of NH3)) with the column densities reported from the observational data.," For these cores, the values obtained from the chemical modeling are in agreement (within a factor of 2 in the case of, and a factor of 4 in the case of ) with the column densities reported from the observational data."305 Finally. the CO fractional abundance in the western/eastern cores is ~107. significantly lower than the CO fractional abundance of the central core. indicating that the relatively low CO abundance in the western/eastern cores does not lead to a substantial destruction ofNoH7.," Finally, the CO fractional abundance in the western/eastern cores is $\sim10^{-5}$, significantly lower than the CO fractional abundance of the central core, indicating that the relatively low CO abundance in the western/eastern cores does not lead to a substantial destruction of."306. Therefore. the fraction of CO that will evaporate from grain mantles plays an important role in determining the fractional abundance ofN3H.. and hence the aabundance ratio.," Therefore, the fraction of CO that will evaporate from grain mantles plays an important role in determining the fractional abundance of, and hence the abundance ratio."307 In Table 5 we show the values of the aand ccolumn densities. together with the aabundanceο” ratio for the two models that are in agreement with the observed values.," In Table \ref{tableres} we show the values of the and column densities, together with the abundance ratio for the two models that are in agreement with the observed values."308 In Fig., In Fig.309 9. (bottom right panel) we present the aabundance ratio as a function of time for the central core and the western/eastern core obtained during phase II., \ref{modelfigs}~ (bottom right panel) we present the abundance ratio as a function of time for the central core and the western/eastern core obtained during phase II.310 The aabundance ratio. observed. toward the central core. around -400-1000. can be reproduced by our chemical model for times ¢=~[0 yr and ¢=(4.5-5.3)x10° yr.," The abundance ratio observed toward the central core, around $\sim$ 400–1000, can be reproduced by our chemical model for times $t\simeq10^4$ yr and $t=$ $\times10^5$ yr."311 For the time range r= 1074.5x10° yr the model produces a higher aabundance ratio. ~4000. slightly above the observed values for the central core.," For the time range $t\simeq10^4$ $4.5\times10^5$ yr the model produces a higher abundance ratio, $\sim4000$, slightly above the observed values for the central core."312 However we adopted the longer age as It is more realistic., However we adopted the longer age as it is more realistic.313 Forthe western/eastern cores. the abundance ratio initially shows high values. then it decreases due to desorption effects and finally reaches a constant value ~200 for typical ages of low-mass YSOs. at around f=10°— yr.," Forthe western/eastern cores, the abundance ratio initially shows high values, then it decreases due to desorption effects and finally reaches a constant value $\sim200$ for typical ages of low-mass YSOs, at around $t\simeq10^{5}-10^{6}$ yr."314 In addition. in Fig.," In addition, in Fig."315 9 (bottom panel) we also show the temperature as a function of time., \ref{modelfigs} (bottom panel) we also show the temperature as a function of time.316 For the central core. when the temperature ts low. the rratio is high (> 107).," For the central core, when the temperature is low, the ratio is high $>10^4$ )."317 Around Τ~21 K there is a clear drop of the rratio due to desorption effects., Around $T\simeq21$ K there is a clear drop of the ratio due to desorption effects.318 As temperature increases. the rrises until it reaches a constant value. ~10°. at around T.— K. For the western and eastern cores the temperature varies from 12 to 25 K. producing small variations on the ratio. and only for times in the range 10°-10' yr the rratio changes significantly. similar to the case of the central core.," As temperature increases, the rises until it reaches a constant value, $\sim10^3$, at around $T\simeq45$ K. For the western and eastern cores the temperature varies from 12 to 25 K, producing small variations on the ratio, and only for times in the range $10^3$ $10^4$ yr the ratio changes significantly, similar to the case of the central core."319 It i5 worth noting that in the model for the central core the visual extinction is Ay=40. but typically the visual extinction in hot cores eembedded in the central core) is around ~100 or even higher.," It is worth noting that in the model for the central core the visual extinction is $A_{\mathrm{V}}\simeq40$, but typically the visual extinction in hot cores embedded in the central core) is around $\sim100$ or even higher."320" In order to evaluate the error in the aabundance introduced by this difference we performed an additional model with a higher density. which gives a visual extinction of A,= 100."," In order to evaluate the error in the abundance introduced by this difference we performed an additional model with a higher density, which gives a visual extinction of $A_{\mathrm{v}}\simeq100$ ."321 In this situation. the rratio is affected by a difference of ~4 3-25 We also note that," In this situation, the ratio is affected by a difference of $\sim3$ –25 .We also note that"322the 6.511 Magellan Daade telescope with the £/2 camera.,the 6.5m Magellan Baade telescope with the f/2 camera.323 All spectra cover a wavelcneth ranee of 1000-0000. With a dispersion of pper pixel and were reduced withCOSALOS. the standard data reduction package for INLAC'S spectra.," All spectra cover a wavelength range of 4000-9000 with a dispersion of per pixel and were reduced with, the standard data reduction package for IMACS spectra."324 To obtain redshifts. we use the IRAF task to cross-correlate spectra with template spectra of four ealaxv types giant elliptical. spiral. E|A. aud eiissiou iue galaxy.," To obtain redshifts, we use the IRAF task to cross-correlate spectra with template spectra of four galaxy types – giant elliptical, spiral, E+A, and emission line galaxy."325 Of the 1122 spectra. we recover redshitts or S57 galaxies. of which 362 are designated cluster uenibers based on a caustic analysis of the cluster iufall region (Gonzalez et al.," Of the 1122 spectra, we recover redshifts for 857 galaxies, of which 362 are designated cluster members based on a caustic analysis of the cluster infall region (Gonzalez et al."326 in prep). snülu to the echuique emploved by Diaferioetal.(2005).," in prep), similar to the technique employed by \citet{diaferio2005}."327. To our catalog of 362 confirmed members. we add another LL ucibers from Barrenaetal. (2002)..," To our catalog of 362 confirmed members, we add another 44 members from \citet{barrena2002}. ."328 Figure 1. shows the distribution of spectroscopic redshifts. with the cluster uembers hiehliehted iu red.," Figure \ref{fig:veldisp} shows the distribution of spectroscopic redshifts, with the cluster members highlighted in red."329 In addition to the Bullet Cluster. we find two prominent redshift peaks at 20.21 and 2~0.35. contaiuiug roughly 90 members each.," In addition to the Bullet Cluster, we find two prominent redshift peaks at $z\sim0.21$ and $z\sim0.35$, containing roughly 90 members each."330 TheObservatory far-intrared and. sub-nuu properties of the backerouncd eroup are studied iu relation to the Bullet Cluster by Rawle et al (2010. in press).," The far-infrared and sub-mm properties of the background group are studied in relation to the Bullet Cluster by Rawle et al (2010, in press)."331 To quantify the total star formation rate of the Bullet Cluster as traced by Ihuninositv. we first identify the MIPS sources that are likely to be cluster members.," To quantify the total star formation rate of the Bullet Cluster as traced by luminosity, we first identify the MIPS sources that are likely to be cluster members."332 We use WEI aud IRAC nuages to exclude galaxies whose colors are indicative of either backerouncd galaxies or active galactic uuclci (ACN)., We use WFI and IRAC images to exclude galaxies whose colors are indicative of either background galaxies or active galactic nuclei (AGN).333 We then refined our sample by spectroscopically tarecting AIIPS sources that are candidate cluster members., We then refined our sample by spectroscopically targeting MIPS sources that are candidate cluster members.334 We start with an initial MIPS catalog of [18 sources down to a flux of [μὴν , We start with an initial MIPS catalog of 418 sources down to a flux of $43\mu$ Jy.335We match the WFI (BYR) and IRAC catalogs and then crossauateh with the MIPS catalog., We match the WFI (BVR) and IRAC catalogs and then cross-match with the MIPS catalog.336 Figure 2. shows the R-[L5] versus D-R colors for all MIPS sources with optical and IRAC counterparts., Figure \ref{fig:select} shows the R-[4.5] versus B-R colors for all MIPS sources with optical and IRAC counterparts.337 Spectroscopically confirmed members from the 2005. 2006. and first part of 2009 INLACS campaigns demonstrated that the star-forming cluster menibers form a tight diagonal locus iu the R-|1.5| vs B-R space.," Spectroscopically confirmed members from the 2005, 2006, and first part of 2009 IMACS campaigns demonstrated that the star-forming cluster members form a tight diagonal locus in the R-[4.5] vs B-R space."338 We define our cluster candidate sample as galaxies within the two solid diagonal lines shown in Figure 2.. and blueward of B-R=3.," We define our cluster candidate sample as galaxies within the two solid diagonal lines shown in Figure \ref{fig:select}, and blueward of B-R=3."339 We then use IRAC colors to ideutiftv AGN via the “AGN ποσο (Lacyctal.2001:Sternet2005).. as shown in Figure 3..," We then use IRAC colors to identify AGN via the “AGN wedge” \citep{lacy2004,stern2005}, as shown in Figure \ref{fig:irac_colors}. ."340 We exclude all galaxies within the AGN wedee from further analyses because their mid-infrared. Iuuinosityv may be dominated by ACN activity rather than bv star formation., We exclude all galaxies within the AGN wedge from further analyses because their mid-infrared luminosity may be dominated by AGN activity rather than by star formation.341" We also exclude X-ray AGN using a catalog of 115 X-ray point sources extracted from Chandra data that cover the central ~20& 207, overlapping the entire MIPS ΕΟΝ."," We also exclude X-ray AGN using a catalog of 145 X-ray point sources extracted from Chandra data that cover the central $\sim20\arcmin\times 20\arcmin$ , overlapping the entire MIPS FOV."342" The catalog contains N-ray sources down to a flux of 2.5«1016 eves au? + in the 0.5-2 keV. band. which is a luninosity of L,.=ς10% eres s tat the redshitt of the Bullet Cluster."," The catalog contains X-ray sources down to a flux of $2.5\times10^{-16}$ ergs $^{-2}$ $^{-1}$ in the 0.5-2 keV band, which is a luminosity of $L_{x}=7\times10^{40}$ ergs $^{-1}$ at the redshift of the Bullet Cluster."343" Áimong the MIPS confirmed cluster members and candidates. we exclude all X-ray point. sources with an N-ray luminosity of £,>10H cress |. Iu addition to IRAC aud N-rav selection of ACN. we also utilize our optical spectroscopy to construct Baldwin-Philips-Terlevich diagrams (BPT:Baldwinctal.1951) as illustrated in Figure L."," Among the MIPS confirmed cluster members and candidates, we exclude all X-ray point sources with an X-ray luminosity of $L_{x}\geq10^{41}$ ergs $^{-1}$ In addition to IRAC and X-ray selection of AGN, we also utilize our optical spectroscopy to construct Baldwin-Phillips-Terlevich diagrams \cite[BPT;][]{baldwin1981} as illustrated in Figure \ref{fig:BPT}."344 The dotted aud dashed lines indicate the boundaries from Ἱνοπ]ονetal.(2006) that separate purely star forming galaxies from Sevferts and LINERS., The dotted and dashed lines indicate the boundaries from \citet{kewley2006} that separate purely star forming galaxies from Seyferts and LINERs.345 Note that many of the MIPS sources are nisse in Figure Lo because one or more of the four required eiission lines could not be measured. often due to a promineut skv cussion line that appears at the same waveleugth as IL) at the Bullet Cluster redshift.," Note that many of the MIPS sources are missing in Figure \ref{fig:BPT} because one or more of the four required emission lines could not be measured, often due to a prominent sky emission line that appears at the same wavelength as $\beta$ at the Bullet Cluster redshift."346 Figure | reveals ouly two MIPS sources classified as a Sevtert or LINER that have not been ideutified as ACN using either IRAC colors or X-ray e1iission., Figure \ref{fig:BPT} reveals only two MIPS sources classified as a Seyfert or LINER that have not been identified as AGN using either IRAC colors or X-ray emission.347 One of these is à. ULIRG (star sviubol). which appears close to the Seyfert/LINER boundary wheu using the /Tlo ratio. and is classified as a LINER using the Πα ratio.," One of these is a ULIRG (star symbol), which appears close to the Seyfert/LINER boundary when using the $\alpha$ ratio, and is classified as a LINER using the $\alpha$ ratio."348 The two confirmed MIPS sources in [OTthe IRAC AGN wedge (Figure 3)) are not shown ou the BPT diagram )ecause they lack the necessary euussion nes., The two confirmed MIPS sources in the IRAC AGN wedge (Figure \ref{fig:irac_colors}) ) are not shown on the BPT diagram because they lack the necessary emission lines.349 However. of the 10 confirmed IRAC. AGN wedge sources that are rot MIPS iuciubers (or bevoud the MIPS FOV). two are Xotted in Figure [. with one classified as à LINER. aud he other as an II dominated galaxy.," However, of the 10 confirmed IRAC AGN wedge sources that are not MIPS members (or beyond the MIPS FOV), two are plotted in Figure \ref{fig:BPT}, with one classified as a LINER, and the other as an HII dominated galaxy."350 Of the three X-rav point sources in Figure 3.. one is illustrated in Figure L asa large square aud appears close o the boundary of III dominated galaxies aud Sevferts.," Of the three X-ray point sources in Figure \ref{fig:irac_colors}, one is illustrated in Figure \ref{fig:BPT} as a large square and appears close to the boundary of HII dominated galaxies and Seyferts."351 Iu total. there are S AGN identified amone the confirmed MIPS sample. using the three methods TRAC colors. N-rav enission. aud optical enission Lue ratios.," In total, there are 8 AGN identified among the confirmed MIPS sample, using the three methods – IRAC colors, X-ray emission, and optical emission line ratios."352 Of these cight. two are identified solely from the BPT diagnostic. one classified as à LINER and the other a Sevtert.," Of these eight, two are identified solely from the BPT diagnostic, one classified as a LINER and the other a Seyfert."353 The LINER is included iu our sample. asstuuine that of its IR fux is powered by star formation (see ULIRC OG).," The LINER is included in our sample, assuming that of its IR flux is powered by star formation (see \\ref{sec:ULIRG}) )."354 The Sevtert. which is excluded from our siuuple. would contribute a neglible fraction to the global SFR.," The Seyfert, which is excluded from our sample, would contribute a neglible fraction to the global SFR."355 Iu addition to isolating the AGN population. Figure 3 also illustrates that our selection of cluster candidates based on R-[1.5] and B-R colors is au effective way to cull interlopers.," In addition to isolating the AGN population, Figure \ref{fig:irac_colors} also illustrates that our selection of cluster candidates based on R-[4.5] and B-R colors is an effective way to cull interlopers."356 The left aud right panels of Figure 3) show the IRAC color distribution before aud after we apply the aud B-R color selection. respectively.," The left and right panels of Figure \ref{fig:irac_colors} show the IRAC color distribution before and after we apply the R-[4.5] and B-R color selection, respectively."357 We show the R-|1.5]model colors ofM82 (a local starburst galaxy) at 2=0.0.3.0.5. and 1 (Devriendtctal.1999:Steru2005) to illustrate that our color selection is effective im removing ealaxies whose IRAC colors are consistent with those of a Ligh redshift starburst (open red circles near οσο1 to 2).," We show the model colors ofM82 (a local starburst galaxy) at z=0,0.3,0.5, and 1 \citep{devriendt1999,stern2005} to illustrate that our color selection is effective in removing galaxies whose IRAC colors are consistent with those of a high redshift starburst (open red circles near $z\sim1$ to $2$ )."358 The right panel of Figure 3. highlights the AGN sources among the MIPS sample. includiug four galaxies iu the IRAC AGN wedge. three N-vay point sources. outside the AGN wedge. aud one Sevfert identified fro the BPT diagram. which is sciui-hidden behind the ULIRG svinbol.," The right panel of Figure \ref{fig:irac_colors} highlights the AGN sources among the MIPS sample, including four galaxies in the IRAC AGN wedge, three X-ray point sources outside the AGN wedge, and one Seyfert identified from the BPT diagram, which is semi-hidden behind the ULIRG symbol."359 Among the mou-ACN MIPS. cluster ποσα». there are two outliers. both of which do not ft in the diagonal star forming sequence (roughly outlined by the LIRCGs and ULIRG shownas star sviubols) nor the locus of passive carly type galaxies near [6|-|L.5|50.2. aud [5.8|-[8.0]~0.2.One outlier has a color of « 0. due toa blending of sources iu theaud Dbauds.," Among the non-AGN MIPS cluster members, there are two outliers, both of which do not fit in the diagonal star forming sequence (roughly outlined by the LIRGs and ULIRG shownas star symbols) nor the locus of passive early type galaxies near $\sim$0.2 and $\sim$ 0.2.One outlier has a color of $<0$ , due to a blending of sources in theand bands."360 It is unclear why the second. outlier has IRAC, It is unclear why the second outlier has IRAC361colours are 1.0«(BG.RR)1.5 which we will use to compare the B band observations of the previous section with the R band observations ciscussed here.,colours are $1.0<(B-R)<1.5$ which we will use to compare the B band observations of the previous section with the R band observations discussed here.362 To test the procedure described in the previous section we have added simulated ares to the raw data (each image is mace up of 5 separate LO min exposures) and then processed the frame as before., To test the procedure described in the previous section we have added simulated arcs to the raw data (each image is made up of 5 separate 10 min exposures) and then processed the frame as before.363 Below we will show that an are of surface brightness 27.5 Rye (28.5-29 Byr) is easily. detectable., Below we will show that an arc of surface brightness 27.5 $\mu$ (28.5-29 $\mu$ ) is easily detectable.364 This is 0.5 to 1.0 magnitudes deeper than the observations described in the previous section., This is 0.5 to 1.0 magnitudes deeper than the observations described in the previous section.365 In Fig., In Fig.366 3 we show the data frame (before smoothing) and the distribution of angles that low surface brightness pixels make with the origin., 3 we show the data frame (before smoothing) and the distribution of angles that low surface brightness pixels make with the origin.367 Looking carefully. the simulated. stream can just be mace out as à diagonal line across the CCD from bottom left to top right.," Looking carefully, the simulated stream can just be made out as a diagonal line across the CCD from bottom left to top right."368 ln the histogram an arc at an angle of 45 degrees is clearly seen at SN=4.4., In the histogram an arc at an angle of 45 degrees is clearly seen at $SN = 4.4$.369 Lt is quite clear that ares at this surface brightness level are not removed by the data processing., It is quite clear that arcs at this surface brightness level are not removed by the data processing.370 Given the strong signallein] obtained for an are at a surface brightness of 27.55 Ry we are confident that our procedure is capable of detecting ares in this data set to an equivalent level of 29 Dj., Given the strong signal obtained for an arc at a surface brightness of 27.5 $\mu$ we are confident that our procedure is capable of detecting arcs in this data set to an equivalent level of 29 $\mu$.371 We have applied the above procedure to the 7 galaxies in our sample., We have applied the above procedure to the 7 galaxies in our sample.372 None of the galaxies have a detection at SUVo3: one has 2«SN3., None of the galaxies have a detection at $SN>3$; one has $2<SN<3$.373 There is no evidence for any low surface brightness arcs., There is no evidence for any low surface brightness arcs.374 Although the original numerical simulations of galaxy lavassment were carried out with a cluster designed to be ike Coma. rather than Virgo. the conclusion was that this was a viable method for creating the excess of dwarf galaxies ound in clusters in general (Moore et al..," Although the original numerical simulations of galaxy harassment were carried out with a cluster designed to be like Coma, rather than Virgo, the conclusion was that this was a viable method for creating the excess of dwarf galaxies found in clusters in general (Moore et al.,"375. 1999)., 1999).376 The clear result from this paper is that we find little or no evidence or any extendedtidal streams. with surface brightnesses Less han 29 Dir. associated with our target. galaxies that would »e the signature of morphological transformation by galaxy iwassment.," The clear result from this paper is that we find little or no evidence for any extendedtidal streams, with surface brightnesses less than 29 $\mu$, associated with our target galaxies that would be the signature of morphological transformation by galaxy harassment."377 There are a number of possible explanations., There are a number of possible explanations.378 The most obvious is that they do not exist. (discussed. in more detail below)., The most obvious is that they do not exist (discussed in more detail below).379 The second is that they are too faint., The second is that they are too faint.380 ‘This could. arise because we have over estimated the stellar uminosity of the progenitor galaxies (the ares are at much ower surface brightnesses) or because the tidal material was ο out of the galaxies a long time ago (more than a ew billion vears) and has since dispersed., This could arise because we have over estimated the stellar luminosity of the progenitor galaxies (the arcs are at much lower surface brightnesses) or because the tidal material was pulled out of the galaxies a long time ago (more than a few billion years) and has since dispersed.381 We cannot rule out either of these two. but we believe that there are other reasons for being skeptical about the harassment. model - he morphological transformation of LSB disc galaxies into dl galaxies: We have previously proposed. that dwarf. galaxies are rather robust objects in the cluster environment. because they have large mass-to-light ratios (Sabatini et al..," We cannot rule out either of these two, but we believe that there are other reasons for being skeptical about the harassment model - the morphological transformation of LSB disc galaxies into dE galaxies: We have previously proposed that dwarf galaxies are rather robust objects in the cluster environment because they have large mass-to-light ratios (Sabatini et al.,"382 2004)., 2004).383 However. in clusters like Virgo they do undergo accelerated evolution (star formation) clue to tidal interactions in the cluster (see for example Lenriksen Byrd. 1996). which eventually exhausts thei gas supply.," However, in clusters like Virgo they do undergo accelerated evolution (star formation) due to tidal interactions in the cluster (see for example Henriksen Byrd, 1996) which eventually exhausts their gas supply."384 “Phus the cluster environment is one that reveals its dark matter halos to us., Thus the cluster environment is one that reveals its dark matter halos to us.385 So. are cluster dwarl galaxies the embers of a once much brighter lire or are they ignited. for the first time in the cluster environment," So, are cluster dwarf galaxies the embers of a once much brighter fire or are they ignited for the first time in the cluster environment ?"386The core radius and central density obtained after a careful decontamination process of Swaters et al. (,The core radius and central density obtained after a careful decontamination process of Swaters et al. (3872000) data (big empty squares) are again in the range of our estimates described previously for the larger body of data. showing in act à somewhat reduced dispersion.,"2000) data (big empty squares) are again in the range of our estimates described previously for the larger body of data, showing in fact a somewhat reduced dispersion."388 This leads us to expect hat i£ sullicient information were available on the svstenis of our larger data set. à more detailed reconstruction of their »uwameters would. similarlv. leac to a reduced. scatter and not show any systematic ollsets.," This leads us to expect that if sufficient information were available on the systems of our larger data set, a more detailed reconstruction of their parameters would similarly lead to a reduced scatter and not show any systematic offsets."389 lig., Fig.390 1-4 show basic properties of the nature of dark ialoes. albeit the scatter seen. which theory must aclelress.," 1-4 show basic properties of the nature of dark haloes, albeit the scatter seen, which theory must address."391 1n a previous work we have shown that a possible physical mechanism capable of vielding the observed scale invariance or the halo central density (Eig., In a previous work we have shown that a possible physical mechanism capable of yielding the observed scale invariance for the halo central density (Fig.392 4) is the selt-interaction of he cold dark particles (Firmani ct al., 4) is the self-interaction of the cold dark particles (Firmani et al.393 2000)., 2000).394 Next. section is dedicated to a review of other processes that have been suggested for producing shallow cores. which however. can »¢ ruled: out because of their disagreement. with the scale invariant behaviour of the halo central density.," Next section is dedicated to a review of other processes that have been suggested for producing shallow cores, which however, can be ruled out because of their disagreement with the scale invariant behaviour of the halo central density."395 The hierarchical CDAL scenario. of structure formation successfully explains most of the main properties of the local and high-redshift universe., The hierarchical CDM scenario of structure formation successfully explains most of the main properties of the local and high-redshift universe.396 However. this scenario does not allow the presence of soft cores in the dark haloes. a point of contlict with the observations.," However, this scenario does not allow the presence of soft cores in the dark haloes, a point of conflict with the observations."397 Here we explore some general alternatives for solving this conllict. without going much into the detailed physical mechanisms.," Here we explore some general alternatives for solving this conflict, without going much into the detailed physical mechanisms."398 A cut-olf in the power spectrum. of Dluetuations at some scale. which avoids the formation of small dense substructures could. be expected το vield. κος cores.," A cut-off in the power spectrum of fluctuations at some scale, which avoids the formation of small dense substructures could be expected to yield soft cores."399 Such a cut-off appears naturally in warm dark. matter scenarios where dark matter particles remain relativistic until relatively late times., Such a cut-off appears naturally in warm dark matter scenarios where dark matter particles remain relativistic until relatively late times.400 In this section our goal is to show that a of the power spectrum of the primeval density [uctuation field fails in producing haloes with both soft. cores and. nearly constant central densities. over the entire mass sampled in Vig.," In this section our goal is to show that a of the power spectrum of the primeval density fluctuation field fails in producing haloes with both soft cores and nearly constant central densities, over the entire mass sampled in Fig.4."401 Reeently Moore. and co-workers (1999) using N-bocky echniques have simulated CDAL haloes taking a lower cut-olf in the primeval power spectrum of the fluctuation ied. as in a warm dark matter universe., Recently Moore and co-workers (1999) using N-body techniques have simulated CDM haloes taking a lower cut-off in the primeval power spectrum of the fluctuation field as in a warm dark matter universe.402 In this case virialized haloes with steep density profiles. are. formed with an asvmptotic central slope pxrl.l, In this case virialized haloes with steep density profiles are formed with an asymptotic central slope $\rho \propto r^{-1.4}$.403 From our volt Of view this result is basically related to the lack of article angular momentum., From our point of view this result is basically related to the lack of particle angular momentum.404 In a similar vein. Huss and co-workers (1999) have investigated several structure formation," In a similar vein, Huss and co-workers (1999) have investigated several structure formation"405xofile in the [Ar 1v] ALLL line compared to the [Ar 1v] ALTLO line.,profile in the [Ar ] $\lambda4711$ line compared to the [Ar ] $\lambda4740$ line.406 For a similar reason. the |Àr iv] A1711 line las a narrower profile than the [Ar tv] Al7LO line for an outwardly increasing density distribution.," For a similar reason, the [Ar ] $\lambda4711$ line has a narrower profile than the [Ar ] $\lambda4740$ line for an outwardly increasing density distribution."407 A comparison of chuission line profiles for various ionic species provides ighter constraiuts oi nebular density structure., A comparison of emission line profiles for various ionic species provides tighter constraints on nebular density structure.408 Using the |S i1] AAG731.6716 line profiles. we calculaed thedeusitv distribution iu velocity space. Noe). for he PNe IC 2501. IC. 1191. and NGC 2[10 (see Sharpee ct al.," Using the [S ] $\lambda\lambda6731,6716$ line profiles, we calculated thedensity distribution in velocity space, $N_{\rm e}(v)$, for the PNe IC 2501, IC 4191, and NGC 2440 (see Sharpee et al."409 2007 for the observation details).," \cite{sharpee07}410 for the observation details)."411 The results are shown in Fie. 7.., The results are shown in Fig. \ref{n_v}.412 An inspection of the figure clearly indicates lat the deusitv variatious in the velocity field are small (less than 1 dex)., An inspection of the figure clearly indicates that the density variations in the velocity field are small (less than 1 dex).413 Uowever. for the same reasons discussed iu Sect.," However, for the same reasons discussed in Sect."414 3.1.2 for Tite). density varia10115 111 he velocity field iav have becu s100thed aud oulv represeut a lower luit to the nebular deusity variations uuless line broadening is dominated by the expansion velocity aud he velocity field Is monotonic along the nebular radius.," \ref{te} for $T_{\rm e}(v)$, density variations in the velocity field may have been smoothed and only represent a lower limit to the nebular density variations unless line broadening is dominated by the expansion velocity and the velocity field is monotonic along the nebular radius."415 Figure 7 shows that the compact PN IC 2501 has an almost homogeneous density structure with au average density of ~1Meca 72, Figure \ref{n_v} shows that the compact PN IC 2501 has an almost homogeneous density structure with an average density of $\sim10^4$ $^{-3}$.416 The slight increase in density at 30kkm/s is uot real considering the large errors iutroduced. by he weakness of the line wines., The slight increase in density at km/s is not real considering the large errors introduced by the weakness of the line wings.417 IC 1191 Is a conipact aiid imregular PN., IC 4191 is a compact and irregular PN.418 Figure 7 shows that it has an asvuuuetrical density structure., Figure \ref{n_v} shows that it has an asymmetrical density structure.419 A condeusation is evideut arouid λαός ou the red side. which las a density abott twice as high as the other nebular regions.," A condensation is evident around km/s on the red side, which has a density about twice as high as the other nebular regions."420 We cannot find a counterpart on the blue side., We cannot find a counterpart on the blue side.421 NGC 2110 clearv shows a negative deusitv eracdieut across the nebula., NGC 2440 clearly shows a negative density gradient across the nebula.422 The deusity at the outer edge of the nebula is lower bvau order of magnitude than at the centre., The density at the outer edge of the nebula is lower byan order of magnitude than at the centre.423 We find that the |8 n] ALOT6 line has a different profile from the [S uf AAGT31.671G lines for all three PNe.," We find that the [S ] $\lambda4076$ line has a different profile from the [S ] $\lambda\lambda6731,6716$ lines for all three PNe."424 However. the [8 n] ALO76 line was not taken into account in the current analysis because its upper level has a higher excitation energy than the [S 11] AAG731.6716. and hence we cannot separate the coutribution of temperature variations to the difference between their profiles.," However, the [S ] $\lambda4076$ line was not taken into account in the current analysis because its upper level has a higher excitation energy than the [S ] $\lambda\lambda6731,6716$, and hence we cannot separate the contribution of temperature variations to the difference between their profiles."425 To demonstrate how line profiles could be used to iieasure nebular pliysical conditions. we considere data for IC £15 to complete a detailed study of line profiles by attempting to match our moctels to observational data.," To demonstrate how line profiles could be used to measure nebular physical conditions, we considered data for IC 418 to complete a detailed study of line profiles by attempting to match our models to observational data."426 The analysis of a larger uuuber of PNe will appear in a separate paper., The analysis of a larger number of PNe will appear in a separate paper.427 Ilüeh-esolution echelle spectra of the PN IC 115 were obtained in 2001 with the Lia Blanco telescope at Cerro Tololo Iuter-American Observatory (Sharpee etal. 2001)).," High-resolution echelle spectra of the PN IC 418 were obtained in 2001 with the 4 m Blanco telescope at Cerro Tololo Inter-American Observatory (Sharpee etal. \cite{sharpee04}) ),"428 covering a wavelength range from 3500 toA., covering a wavelength range from 3500 to.429 IC. [18 is a low-excitaion PN [F(|O ΠΠ A5007)/FUL3)~ 2] an has an approximatively-splerical shell., IC 418 is a low-excitation PN $F$ ] $\lambda5007)/F({\rm H}\beta)\sim2$ ] and has an approximatively-spherical shell.430" The observations were carried out using a~1"" slit width. vieldiug a spectra resolution of ~30.000."," The observations were carried out using a $\sim1''$ slit width, yielding a spectral resolution of $\sim30,000$."431 The spectrograph slit was aligue north-south., The spectrograph slit was aligned north-south.432 The central star was placed along a line perpendicular to the spectrograph slit so that the slit contro Was roughiv müdway betweenthe ceutral star iux outer edges of he nebula., The central star was placed along a line perpendicular to the spectrograph slit so that the slit centre was roughly midway betweenthe central star and outer edges of the nebula.433 A large nunuber of ORL ux CEL profiles arc well resolved., A large number of ORL and CEL profiles are well resolved.434 yung et al. (1991)), Hyung et al. \cite{hyung94}) )435 constructed a photoionization model of IC. £18 that cau interpret the UV data fairly well., constructed a photoionization model of IC 418 that can interpret the UV data fairly well.436 For our model. we initially took the sic parameters used by νο ct al.," For our model, we initially took the basic parameters used by Hyung et al."437 1991 (disance. properties of the central star. chemical composition. nebular ecometrical structure. ete).," \cite{hyung94}438 (distance, properties of the central star, chemical composition, nebular geometrical structure, etc.),"439 which were then adjusted slightlv to achieve an optimal ft to optical observations., which were then adjusted slightly to achieve an optimal fit to optical observations.440 The density distribution of IC [18 is a topic of ongoing debate., The density distribution of IC 418 is a topic of ongoing debate.441 Tsung ct ab (1991)), Hyung et al. \cite{hyung94}) )442 presented that he nebula may consist of two shells an iuuer hieh-deusitv one aud an outer lower-density shell.," presented that the nebula may consist of two shells, an inner high-density one and an outer lower-density shell."443 À conflicting conclusion was obtained bw Ctesicki ct al. (1996)), A conflicting conclusion was obtained by Gesicki et al. \cite{gesicki96}) )444 who fouud that the density iereases smoothly outwards by fitting the profiles for U3. |O 11] AS007. and [N 11] AG5sL.," who found that the density increases smoothly outwards by fitting the surface-brightness profiles for $\beta$ , [O ] $\lambda5007$ , and [N ] $\lambda6584$ ."445 Our model assuined a single shell with a homogeneous density of qp=10! ., Our model assumed a single shell with a homogeneous density of $N_{\rm H}=10^4$ $^{-3}$ .446 Below we show that this, Below we show that this447he amaxima of the cleveu-vear evcles in the Sun (Oliverctal.1998:Zaqirashviliet2010). and displaving a periodicity of 150.—160 davs.,"the maxima of the eleven-year cycles in the Sun \citep{Oliveretal98,Zaqarashvilietal10} and displaying a periodicity of $150-160$ days."448 According o Lou(2000) and Zaqarashvilietal.(2010).. Bossby-tvpe magnuetolhvdrodvuzanuie waves. excited iu he subphotospheric lavers or at the interface between he convection zone aud the radiative interior of the Sun. could account for such a short-term periodicity bv nodulating the emergeuce of iiagnetic fiux.," According to \citet{Lou00} and \citet{Zaqarashvilietal10}, Rossby-type magnetohydrodynamic waves, excited in the subphotospheric layers or at the interface between the convection zone and the radiative interior of the Sun, could account for such a short-term periodicity by modulating the emergence of magnetic flux."449 Since the xeriodof the wave is proportional to the rotationperiod of the star. we expect a period of about one mouth iu he case of CoRoT-2 that is a C-type star rotating five ines faster than the Sun.," Since the period of the wave is proportional to the rotation period of the star, we expect a period of about one month in the case of CoRoT-2 that is a G-type star rotating five times faster than the Sun."450 Ax alternative explanation for the modulation of tle spotted area in CoRoT-2 considers it as a signature of SPMI in the photosphere of the star period., An alternative explanation for the modulation of the spotted area in CoRoT-2 considers it as a signature of SPMI in the photosphere of the star .451" The svnodic period P4, is the time interval between two successivo passages of the same meridian across the subplanctary longitude in a rotating star.", The synodic period $P_{\rm syn}$ is the time interval between two successive passages of the same meridian across the subplanetary longitude in a rotating star.452" It is equal to 2.89 davs in the case of CoRoT-2. assmuine a mean rotation period 4,4,=1.5221 davs because it is given by: Pt=|P.3P."," It is equal to $2.89$ days in the case of CoRoT-2, assuming a mean rotation period $P_{\rm rot} = 4.5221$ days because it is given by: $P_{\rm syn}^{-1} = | P_{\rm rot}^{-1} - P_{\rm orb}^{-1} | $."453" According to this conjecture. the passage of the planet over an active region inav trigecr the emergence of maenetic flux tubes when their iuteusitv is already close to the threshold for he onset of the Ὀποναπον instability (Acheson197s, 1979)."," According to this conjecture, the passage of the planet over an active region may trigger the emergence of magnetic flux tubes when their intensity is already close to the threshold for the onset of the buoyancy instability \citep{Acheson78, Acheson79}."454. Ten passages are to clapse before the maguctic Ποια intensity reaches again the threshold needed for he planetaryanduced perturbation to be effective., Ten passages are to elapse before the magnetic field intensity reaches again the threshold needed for the planetary-induced perturbation to be effective.455 Spots occulted by the planet during transit can ο detected through the characteristic light Japs xoduced along the transit light curve when the plauct’s dise passes over them., Spots occulted by the planet during transit can be detected through the characteristic light bumps produced along the transit light curve when the planet's disc passes over them.456 Silva-Valio&Lanza(2011) investigated he variation of the total area. and found a periodicity of 17.7£2.3 davs., \citet{Silva-ValioLanza11} investigated the variation of the total area and found a periodicity of $17.7 \pm 2.3$ days.457 Although very close to teu orbital periods of the planet. it is uulikely to be an effect of the aliasing due to the periodic time sampling of the observations.," Although very close to ten orbital periods of the planet, it is unlikely to be an effect of the aliasing due to the periodic time sampling of the observations."458 Iu that case. we should observe other periodicities at. sav 5 or 15 orbital periods which is not the case (secFig.10iuSilva-Valio&Lanza 2011).," In that case, we should observe other periodicities at, say, 5 or 15 orbital periods which is not the case \citep[see Fig.~10 in][]{Silva-ValioLanza11}."459. Moreover. computing the veriodoeramsC» of 10 000 random permutations of the time series. a perlodicity of 17.7 davs occurs oulv iu ~3 percent of the cases.," Moreover, computing the periodograms of 10 000 random permutations of the time series, a periodicity of 17.7 days occurs only in $\sim 3$ percent of the cases."460 The periodicity corresponds to 6 svuodic periods when the rotation period of the mean latitude of the strip occulted by the planet is considered. 1.0. 1.15 davs (Silva-Valio&Lanza2011).," The periodicity corresponds to $\sim 6$ synodic periods when the rotation period of the mean latitude of the strip occulted by the planet is considered, i.e., 4.48 days \citep{Silva-ValioLanza11}."461. The occulted spots are a subset of all the spots present on the star and are ocated at a latitude between 5? and ~207., The occulted spots are a subset of all the spots present on the star and are located at a latitude between $\sim 5^{\circ}$ and $\sim 20^{\circ}$.462 On the other haud. the modulation of the leht curve outside rausit is produced by starspots at all latitudes.," On the other hand, the modulation of the light curve outside transit is produced by starspots at all latitudes."463" We expect that spots cau be formed in a latitude rauge nore extended than in the Sun because CoRoT-2 is at least 20 times more active than our star (ο,ο,Strass-ucier2009).", We expect that spots can be formed in a latitude range more extended than in the Sun because CoRoT-2 is at least 20 times more active than our star \citep[e.g. ][]{Strassmeier09}.464. Therefore. the periodicity of the starspot area of the baud occulted during transits is diluted when we consider the starspot area derived from the out-of-trausit light curve.," Therefore, the periodicity of the starspot area of the band occulted during transits is diluted when we consider the starspot area derived from the out-of-transit light curve."465 As a matter of fact. a peak. corresponding to the second harmonic of the 17.7 day period is present in the periodogram of the out-oftransit spotted area although its height is remarkably lower thau that of the main peak at 29 days because of the effects of the dilution (cf.Fig.7ofLanzactal.20092 ).," As a matter of fact, a peak corresponding to the second harmonic of the 17.7 day period is present in the periodogram of the out-of-transit spotted area although its height is remarkably lower than that of the main peak at 29 days because of the effects of the dilution \citep[cf. Fig.~7 of ][]{Lanzaetal09a}."466. A possible interpretation of the 17.7 day periodicity that corresponds to ~6 svnoclic periods. is that the passage of the planet over the low-latitude active reeions inthe occulted baud iuduces the same phenomenon as observed iu the variation of the fulbdise area.," A possible interpretation of the 17.7 day periodicity that corresponds to $\sim 6$ synodic periods, is that the passage of the planet over the low-latitude active regions in the occulted band induces the same phenomenon as observed in the variation of the full-disc area."467 However. its cadence is shorter because the magnetic field is amplified more rapidly at low latitudes.," However, its cadence is shorter because the magnetic field is amplified more rapidly at low latitudes."468 This interpretation assimes that there are two latitudes where the conjectured mechamisi is operating. one in the occulted band. the other in the latitude rauge uot occulted bv the planet.," This interpretation assumes that there are two latitudes where the conjectured mechanism is operating, one in the occulted band, the other in the latitude range not occulted by the planet."469 The localization of the trigeeringao in two latitude bands is required by the absence of intermediate periodicities in the area modulatiou., The localization of the triggering in two latitude bands is required by the absence of intermediate periodicities in the area modulation.470§ Moreover. these periodicities appear only when the total spotted area. 1.6.. integrated over the longitudo. is considered.," Moreover, these periodicities appear only when the total spotted area, i.e., integrated over the longitude, is considered."471 Therefore. the modulation is not a property of some specific active region. but of all the reeious in a eiven latitude interval.," Therefore, the modulation is not a property of some specific active region, but of all the regions in a given latitude interval."472" Finally, Paganoctal.(2000) considered the variance of the dus of CoRoT-2 in the CoRoT white passband vs. the orbital phase."," Finally, \citet{Paganoetal09} considered the variance of the flux of CoRoT-2 in the CoRoT white passband vs. the orbital phase."473 The variance was calculated using the individual 32 s exposures biuned vs. the orbital phase., The variance was calculated using the individual 32 s exposures binned vs. the orbital phase.474 Cousideriugthe first 75 of davsthe helt curve when the jitter effects were minimal and a bin size of 0.05 inphase. they found that the variance of the flux varied regularly with the orbital phase reaching a maximum just innmediatelv before the trausit and a," Considering the first 75 days of the light curve when the jitter effects were minimal and a bin size of 0.05 inphase, they found that the variance of the flux varied regularly with the orbital phase reaching a maximum just immediately before the transit and a"475"amplitude case> (TG53B1e8) in which the field grows from |B]~3x107? to ~0.02 at radius 0.65@, by runs end.",amplitude case (TG53B1e8) in which the field grows from $|B| \sim 3\times10^{-5}$ to $\sim 0.02$ at radius $\varpi_E$ by runs end.476" A confirmation of the convergence of these results is provided by the high-resolution run TG53B100HR which falls between the B10 and B500 results, and tracks closely the corresponding energy and field amplitudes at lower resolution."," A confirmation of the convergence of these results is provided by the high-resolution run TG53B100HR which falls between the B10 and B500 results, and tracks closely the corresponding energy and field amplitudes at lower resolution."477" The only difference between the high and lower resolution curves is the growth phase in the high resolution case is triggered about one dynamical time earlier (approximately a shift in time), but saturates at the same mean field amplitude."," The only difference between the high and lower resolution curves is the growth phase in the high resolution case is triggered about one dynamical time earlier (approximately a shift in time), but saturates at the same mean field amplitude."478 The exponential growth of the magnetic energy around a time t~11 suggests the onset of a powerful magnetic field amplification mechanism., The exponential growth of the magnetic energy around a time $t \sim 11$ suggests the onset of a powerful magnetic field amplification mechanism.479 The delayed onset and steep exponential growth particularly favor axisymmetric modes of the MRI., The delayed onset and steep exponential growth particularly favor axisymmetric modes of the MRI.480" These modes have the shortest growth times, but require a poloidal field component to act upon."," These modes have the shortest growth times, but require a poloidal field component to act upon."481" For the initially toroidal configurations, significant poloidal fields are not present until the bar mode begins redistributing material within the star, explaining the delayed onset."," For the initially toroidal configurations, significant poloidal fields are not present until the bar mode begins redistributing material within the star, explaining the delayed onset."482" Furthermore, it is clear from Figures 4 and 5 that the magnetic field saturates"," Furthermore, it is clear from Figures \ref{fig:fieldenergy} and \ref{fig:shellamp} that the magnetic field saturates"483"In order to test the ability of our elliptical shapelet pipeline to deal with inaccurate input catalogs, we biased the input ellipticity by 2096 toward circular objects, and measured the residual bias left in the recovered ellipticity 7°.","In order to test the ability of our elliptical shapelet pipeline to deal with inaccurate input catalogs, we biased the input ellipticity by $20\%$ toward circular objects, and measured the residual bias left in the recovered ellipticity $\tilde{\gamma}^{(el)}$."484" As shown by Fig. 4,,"," As shown by Fig. \ref{fig:elliptical-nmax},"485" this test produced conceptually similar results as with circular shapelets: for small Sérrsic indices, we could recover the true ellipticity without any significant bias, while the estimates degrade significantly as we approach n; 4."," this test produced conceptually similar results as with circular shapelets: for small Sérrsic indices, we could recover the true ellipticity without any significant bias, while the estimates degrade significantly as we approach $n_s = 4$ ."486" The situation is improved when 7,4; is raised, because the transformations done during the focusing step take all available orders into account, in contrast to the simpler description underlying the construction of 5? in(14).."," The situation is improved when $n_{max}$ is raised, because the transformations done during the focusing step take all available orders into account, in contrast to the simpler description underlying the construction of $\tilde\gamma^{(n2)}$ in."487" With n4,=12 (bottom panel of Fig. 4)),"," With $n_{max}=12$ (bottom panel of Fig. \ref{fig:elliptical-nmax}) ),"488" galaxies with n,=0.5 and 1 have shear estimates without bias.", galaxies with $n_s=0.5$ and 1 have shear estimates without bias.489" When raising n, beyond that, the bias is at first positive before it becomes negative."," When raising $n_s$ beyond that, the bias is at first positive before it becomes negative."490" Investigating this feature more closely, we find this estimator to have an oscillatory tendency once the bias sets in."," Investigating this feature more closely, we find this estimator to have an oscillatory tendency once the bias sets in."491" In summary, it appears that the pipeline is unable to fully correct inaccuracies in the SExtractor catalog for galaxies showing a profile close to the De Vaucouleurs one."," In summary, it appears that the pipeline is unable to fully correct inaccuracies in the SExtractor catalog for galaxies showing a profile close to the De Vaucouleurs one."492" In more realistic simulations or observational data, the galactic shapes are recorded after convolution with the PSF, pixelation by the CCD, and degraded by pixel noise."," In more realistic simulations or observational data, the galactic shapes are recorded after convolution with the PSF, pixelation by the CCD, and degraded by pixel noise."493 We now discuss the impact of these effects on shear estimation with shapelets., We now discuss the impact of these effects on shear estimation with shapelets.494" Clearly, a convolution creates shallower profiles which can be better described by shapelet models."," Clearly, a convolution creates shallower profiles which can be better described by shapelet models."495" Therefore, the typical goodness-of-fit values, in particular for steeper profiles, are considerable lower than in the unconvolved case."," Therefore, the typical goodness-of-fit values, in particular for steeper profiles, are considerable lower than in the unconvolved case."496" If the PSF shape is perfectly described by its shapelet model, one can exactly undo a convolution in shapelet space."," If the PSF shape is perfectly described by its shapelet model, one can exactly undo a convolution in shapelet space."497" In such a case, the shape obtained by deconvolving a PSF-convolved galaxy model must approximate the true, unconvolved shape G’ better than its direct model G."," In such a case, the shape obtained by deconvolving a PSF-convolved galaxy model must approximate the true, unconvolved shape $G^\prime$ better than its direct model $\tilde{G}^\prime$."498" To verify this new hypothesis, we convolved Sérrsic-type galaxies in pixel space with PSF shapes P obtained from shapeletmodels*,, For circular shapelets, C is modeled with shapelets and explicitly deconvolved from P in shapelet space, while for the elliptical shapelets we obtain the unconvolved shape by convolving the model with the PSF and fitting the outcome to the image data."," To verify this new hypothesis, we convolved Sérrsic-type galaxies in pixel space with PSF shapes $P$ obtained from shapelet, For circular shapelets, $C$ is modeled with shapelets and explicitly deconvolved from $\tilde{P}$ in shapelet space, while for the elliptical shapelets we obtain the unconvolved shape by convolving the model with the PSF and fitting the outcome to the image data."499 In Fig., In Fig.500 5 we compare the bias of the shear estimates from unconvolved galaxies images and from the same set of galaxies after convolution with a Gaussian PSF with a FWHM of 5, \ref{fig:bias-convolution} we compare the bias of the shear estimates from unconvolved galaxies images and from the same set of galaxies after convolution with a Gaussian PSF with a FWHM of 5501"The stars in our normalization sample are selected from the CNS4 by their absolute visual magnitudes. according to the distribution of absolute magnitudes of the CADIS disk stars (6.5<M,1315). Fig.","The stars in our normalization sample are selected from the CNS4 by their absolute visual magnitudes, according to the distribution of absolute magnitudes of the CADIS disk stars $6.5\leq M_v \leq 14.5$ Fig."502 5 shows the resulting density distribution of the disk stars in the two CADIS fields., \ref{Disk} shows the resulting density distribution of the disk stars in the two CADIS fields.503 The solid line represents a fit with a superposition of two exponentials. the dotted line is the fit for the thin disk component (the first seven data points).," The solid line represents a fit with a superposition of two exponentials, the dotted line is the fit for the thin disk component (the first seven data points)."504 Obviously a single exponential is not a good It was suggested to fit the thin disk with a secans hyperbolicus — the exponential is unphysical in that sense. that it is not continuously differentiable at ~=0.," Obviously a single exponential is not a good It was suggested to fit the thin disk with a secans hyperbolicus – the exponential is unphysical in that sense, that it is not continuously differentiable at $z=0$."505 A squared secans hyperbolicus (which represents a self-gravitating isothermal disk) can be proved not to fit the data very well — indicating that the stellar disk is in no way isothermal (the velocity dispersion depends on the spectral The fits for the three functions under consideration: areshown in Fig. 6..," A squared secans hyperbolicus (which represents a self-gravitating isothermal disk) can be proved not to fit the data very well – indicating that the stellar disk is in no way isothermal (the velocity dispersion depends on the spectral The fits for the three functions under consideration: areshown in Fig. \ref{sech},"506 the corresponding parameters are given in Table 2.., the corresponding parameters are given in Table \ref{scales}.507 As the secans hyperbolicus is a sum of two exponentials zy Is not really a scaleheight. but has to be compared to 74 by multiplying it with aresech+zm 1.65715: hj=ze:1.65715 ," As the secans hyperbolicus is a sum of two exponentials $z_0$ is not really a scaleheight, but has to be compared to $h_1$ by multiplying it with $\frac{1}{{\rm e}}\approx 1.65745$ : $h_1'= z_0\cdot5081.65745$ "509"frou, O97.",from O97.510 This means that the reprocessing fraction is simular in all the observations. inuplvius that the amount of reprocessing does not depend strouely on δω.," This means that the reprocessing fraction is similar in all the observations, implying that the amount of reprocessing does not depend strongly on $\Phi _{35}$."511 The best-fit equivalent blackbody radii are cousisteut with, The best-fit equivalent blackbody radii are consistent with512disks: we take Qa;21.,disks; we take $Q_M\gsim 1$.513" Right at the mILR (r~10""). the shortwaveleueth is A4=|X»SΓονμι)~300 oc. While the long waveleneth(3 Ap, would ]© Tufte: at he reflection point of the Qa;,-l ΙΟ (r6c9)30""). they )conie equal Ag=Ap|4o12)/(Gyiy)~600 c."," Right at the mILR $r\sim 10''$ ), the shortwavelength is $\lambda_S=(C_A^2+C_S^2+v_c^2)/(G\mu_{0})\sim 300$ pc, while the long wavelength $\lambda_L$ would be infinite; at the reflection point of the $Q_M$ -barrier $r\sim 30''$ ), they become equal $\lambda_S=\lambda_L=2(C_A^2+C_S^2+v_c^2)/(G\mu_{0})\sim 600$ pc."514" For a mean radial wave2(03 lenethscale of Aa;~15! pe. he damping timescale would be 7,~19,SO.Late s. The timescale τι, for a FMDW to 2oMy from the mILR o Qay-barricr and back to the mILR is estimated o be τιA108 s This, a sizable yaction of negative augular uomientun carried by loug-trailing FPMDWs excite aud sustained at the mILR would ive been deposited in the eas disk as a weaker short-railing FMDW reurus to the uILR."," For a mean radial wave lengthscale of $\lambda_M\sim 450$ pc, the damping timescale would be $\tau_c\sim\lambda_M^2/\nu_c\cong 6\times 10^{15}$ s. The timescale $\tau_w$ for a FMDW to go from the mILR to $Q_M$ -barrier and back to the mILR is estimated to be $\tau_w\sim 3\times 10^{15}$ s. Thus, a sizable fraction of negative angular momentum carried by long-trailing FMDWs excited and sustained at the mILR would have been deposited in the gas disk as a weaker short-trailing FMDW returns to the mILR."515 As the disk aneular momeutuni is reduced. easons naterials with frozen-in magneic flux from outside tend to accumulate just inside the WILR to form a circmmd«§clear starburst “rine”.," As the disk angular momentum is reduced, gas materials with frozen-in magnetic flux from outside tend to accumulate just inside the mILR to form a circumnuclear starburst “ring""."516 An estimate of uct mass inflow raο AL involves several uncertain aspects because the maguitude of excited FAIDWs relaes to the strength of oF and the wave diuping distribution depends ou he effective viscosity as well as sole nonlinear effects., An estimate of net mass inflow rate $\dot M$ involves several uncertain aspects because the magnitude of excited FMDWs relates to the strength of $\phi^E$ and the wave damping distribution depends on the effective viscosity as well as some nonlinear effects.517" By taking a e,|~LOkins l. anda adafemO|lr-:Opn:)E lan (D). an upper linüt of AP av be estimated as kinsMz210Movrb"," By taking a $|v_r|\sim 10\hbox{ km s}^{-1}$ and a $|\omega-m\Omega|r\sim 30\hbox{ km s}^{-1}$ in (4), an upper limit of $\dot M$ may be estimated as $\dot M\lsim 10M_{\odot}\hbox{yr}^{-1}$."518 For a stronger o. nonlinear wave aud ¢aupiue effects may become mniportaut auk this upper Iii nav be raised.," For a stronger $\phi^E$ , nonlinear wave and damping effects may become important and this upper limit may be raised."519 5.5a, 8.5cm520In order to explain several discrepancies between models and near-infrared observations of late IX and AI giants and supereiants. also showing water vapor. Tsuji and collaborators (seeforexampleTsujietal.1997.1998:200Q0a.b:Yamennra1999) have introduced the idea of a stationary. warm envelope situated at a clistance of a few stellar radii above the photosphere but interior to the cool. expanding cireumstellar-shell.,"In order to explain several discrepancies between models and near-infrared observations of late K and M giants and supergiants, also showing water vapor, Tsuji and collaborators \citep[see for521example][]{tsuji_1997,tsuji_1998,tsuji_ny,tsuji_2000,yam_99} have introduced the idea of a stationary, warm envelope situated at a distance of a few stellar radii above the photosphere but interior to the cool, expanding circumstellar-shell."522 This previously undetected envelope (called (he MOLsphere) is considered to contain water vapor al temperatures of 1000. 2000Ix. (Tsujietal.1997).. resulting in non-photospherie signatures in IR. spectra of M giants.," This previously undetected envelope (called the MOLsphere) is considered to contain water vapor at temperatures of $1000$ $2000\,\mbox{K}$ \citep{tsuji_1997}, resulting in non-photospheric signatures in IR spectra of M giants."523 The envelope was neither theoretically. predicted nor has it as vel received a theoretical explanation. but the water lines and bands seem (o be a common feature of M superegiants and M giants (Tsujietal.1998:Matsuura1999).," The envelope was neither theoretically predicted nor has it as yet received a theoretical explanation, but the water lines and bands seem to be a common feature of M supergiants and M giants \citep{tsuji_1998,matsuura}."524.. Numerous pieces of evidence have been presented in favor of this idea (see[orexampleTsujiTsuji 2000b).," Numerous pieces of evidence have been presented in favor of this idea \citep[see for example][]{tsuji_1997, tsuji_1998,tsuji_2000}."525. IIere. we will. however. argue (hat the water-vapor lines detected in the observations of a Boo are not Irom a MOLsphere but of photospheric origin. a possibility anticipated by Tsuji(1935).," Here, we will, however, argue that the water-vapor lines detected in the observations of $\alpha$ Boo are not from a MOLsphere but of photospheric origin, a possibility anticipated by \citet{tsuji:88}."526. Very little work has been done on resolved molecular lines in (he mid-infrared region., Very little work has been done on resolved molecular lines in the mid-infrared region.527 Our analvsis is based on spectra of Arcturus obtained with the spectrograph. a unique. high-resolution. mid-infrared spectrograph.," Our analysis is based on spectra of Arcturus obtained with the spectrograph, a unique, high-resolution, mid-infrared spectrograph."528 This is the [ist time. to our knowledge. this wavelength region has been observed at high-resolution with high sensitivity.," This is the first time, to our knowledge, this wavelength region has been observed at high-resolution with high sensitivity."529 In Section 2 we present the observations and in Sect., In Section 2 we present the observations and in Sect.530 3 we identilv the observed features ancl present our line lists which are used in the generation of svnthetic spectra. a process described in sect.," 3 we identify the observed features and present our line lists which are used in the generation of synthetic spectra, a process described in Sect."531 4., 4.532 Section 5 we discussour findings., Section 5 we discussour findings.533 We observed Arcturus at 806.3—821.4em. 4. 883.8—901.6em. +. and 903.6—922.7em. on 2001 February 2. 4. ancl 3. respectively. with the Texas Echelon-Cross-Echelle Spectrograph(TEXES: Lacy οἱ al.," We observed Arcturus at $806.3-821.4 \,\mbox{cm$ $}$, $883.8-901.6 \,\mbox{cm$ $}$, and $903.6-922.7534\,\mbox{cm$ $}$ on 2001 February 2, 4, and 3, respectively, with the Texas Echelon-Cross-Echelle Spectrograph; Lacy et al."535 mounted on the 3 meter NASA Infrared Telescope Facility (RTI)., \nocite{texes} mounted on the 3 meter NASA Infrared Telescope Facility ).536 is a ground-based prototvpe ofEXES. a mid-intrared (350—1800em !) spectrograph designed lor use on (the Stratospheric Observatory For Infrared Astronomy).," is a ground-based prototype of, a mid-infrared $350-1800\,\mbox{cm$ $}$ ) spectrograph designed for use on (the Stratospheric Observatory For Infrared Astronomy)."537 We observed with a resolving power of ολοzz80.000. » being the wavenumber.," We observed with a resolving power of $\tilde{\nu}/\Delta\tilde{\nu}\approx 80,000$, $\tilde{\nu}$ being the wavenumber."538 A sample of the observed spectrum al 815—5S1Tem is shown in Figure 1..," A sample of the observed spectrum at $815-817\,\mbox{cm$ $}$ is shown in Figure \ref{figsample}. ."539 The lines identified by comparing with the sunspot spectrum (Wallaceetal.1994). are pure rotation lines of water, The lines identified by comparing with the sunspot spectrum \citep{solspectrum} are pure rotation lines of water540 (Freeman&Rodecrs1975:FreemanNorris1981," \citep{FR:75,FN:81,MP:81}. \citealt{SK:96};"541).. (Suutzeff&Ixraft1996: AMevlan Ikuta&Arimoto20003). (Norrisctal.1997)., \citealt{MH:97} \citealt{IA:00}) \citep{Metal:01}. \citep{Netal:97}.542. (~7 D) Hughe, $\sim 7$ $^{-1}$ \citep{FF:82}.543s&Wallerstein2000:IDlker2000:Leeetal.1999)). (Norris&DaCosta1995:Vauture.Waller1991). (Ziuneckeretal.1988:Freeman19," $\Delta\tau544= 3-5$ \citealt{HW:00,HR:00,Letal:99}) \citep{NDC:95,VWB:94}. \citealp{Netal:96,Setal:00}) \citep{Zetal:88,F:93}. \citet{H:96}."54593).. Hiis(1996)...," $M/L_V = 3$ \citet{PM:93},"546a wider field survey on the cluster. uncovering 42 variables.,"a wider field survey on the cluster, uncovering 42 variables."547 We present a new extensive variable star catalogue. a natural byproduct. of a photometric project (ο search for transiting “Lot Jupiter planets in 47 Tuc.," We present a new extensive variable star catalogue, a natural byproduct of a photometric project to search for transiting 'Hot Jupiter' planets in 47 Tuc."548 The results of the planetary search will be the subject of a separate paper., The results of the planetary search will be the subject of a separate paper.549 The total catalogue comprises 28 Eclipsing Binary svstems. 20 long period variables and 41 Small Magellanic Cloud (SAIC) RR Lyraes. four Halo RR Lyraes. (wo Cepheids and four apparent 6 Scuti stars. and one anomalous short period Small Magellanic Cloud (SAIC) star.," The total catalogue comprises 28 Eclipsing Binary systems, 20 long period variables and 41 Small Magellanic Cloud (SMC) RR Lyraes, four Halo RR Lyraes, two Cepheids and four apparent $\delta$ Scuti stars, and one anomalous short period Small Magellanic Cloud (SMC) star."550 Discrimination between Cluster/SAIC memberships is achieved using the location of the variable on the cluster Colour Magnitude Diagram., Discrimination between cluster/SMC memberships is achieved using the location of the variable on the cluster Colour Magnitude Diagram.551 Our large number of new discoveries is due primarily to the very luge field of view (52x52/) of our survey andl also its photometric depth., Our large number of new discoveries is due primarily to the very large field of view $\times$ $'$ ) of our survey and also its photometric depth.552 In this paper. we present (he pliase-wrapped. ΕΙ lighteurves. preliminary investigations into the detected. variable stars. ancl a description of our photometric. astrometric and lighteurve database.," In this paper, we present the phase-wrapped V+R lightcurves, preliminary investigations into the detected variable stars, and a description of our photometric, astrometric and lightcurve database."553 Our survev covers a larger area (han any previous search. aud extends to deeper photometry (han presented by IxXaluzuy.etal.(1998).," Our survey covers a larger area than any previous search, and extends to deeper photometry than presented by \citet{Kal98}."554. We recover 31 of Ixaluznys stars and discover a further 69 variables., We recover 31 of Kaluzny's stars and discover a further 69 variables.555 The unrecovered. variables either lie between. CCDs. or within regions of no data caused by telescope offsets.," The unrecovered variables either lie between CCDs, or within regions of no data caused by telescope offsets."556 The cluster core cannot be easily imaged by ground-based telescopes due to the extreme crowding., The cluster core cannot be easily imaged by ground-based telescopes due to the extreme crowding.557 On our 300s exposures (V+R). the inner 6’ of 47 Tue is saturated.," On our 300s exposures (V+R), the inner $'$ of 47 Tuc is saturated."558 The cluster field is located at a high galactic latitude (=305.9 deg. b=—44.9 deg). providing low loreeround contamination bv the Milky Way and low reddening.," The cluster field is located at a high galactic latitude (l=305.9 deg, $-$ 44.9 deg), providing low foreground contamination by the Milky Way and low reddening."559 The field is significantly contaminated by background stars [rom the Small Magellanic Cloud., The field is significantly contaminated by background stars from the Small Magellanic Cloud.560 Our field of view extends to 76056 of the 47 Tue tidal radius., Our field of view extends to $\sim$ $\%$ of the 47 Tuc tidal radius.561 Study of the SAIC RRLviae stars. as standard candles. presents an opportunity (o investigate the distance to the SAIC from a location some seven degrees NW of the centre.," Study of the SMC RRLyrae stars, as standard candles, presents an opportunity to investigate the distance to the SMC from a location some seven degrees NW of the centre."562 Contact eclipsing binaries (EcDB) are very. useful as distance indicators., Contact eclipsing binaries (EcB) are very useful as distance indicators.563 Observing any, Observing any564the dayside spectra in Section 4.1.,the dayside spectra in Section 4.1.565" Also, systematic errors inherent in the observations themselves could be a strong error source."," Also, systematic errors inherent in the observations themselves could be a strong error source."566 These two facts may cause the planet-star flux ratio to be dependent on the data reduction process., These two facts may cause the planet-star flux ratio to be dependent on the data reduction process.567 Non concurrent observations for different wavelengths could cause large variations in the measured flux between observations if there is any significant temporal variability in the atmosphere of HD 189733b., Non concurrent observations for different wavelengths could cause large variations in the measured flux between observations if there is any significant temporal variability in the atmosphere of HD 189733b.568" Therefore, all these sources of potential error could lead to substantial inconsistencies between the datasets taken from different studies, instruments, and observation times."," Therefore, all these sources of potential error could lead to substantial inconsistencies between the datasets taken from different studies, instruments, and observation times."569" Although we do not intend to revisit the decorrelation procedures with this analysis, we aim to show the limitations of the spectra available at the present epoch under the assumption that the measurements are an accurate representation of the mean planetary flux."," Although we do not intend to revisit the decorrelation procedures with this analysis, we aim to show the limitations of the spectra available at the present epoch under the assumption that the measurements are an accurate representation of the mean planetary flux."570" The molecular composition of the dayside atmosphere of HD 189733b is generally accepted to be a mixture ofH2O, CO», CO, and CH,4."," The molecular composition of the dayside atmosphere of HD 189733b is generally accepted to be a mixture of$_{2}$ O, $_{2}$, CO, and $_{4}$."571" These molecules have been discovered based on various spectroscopic measurements, listed in Section 3."," These molecules have been discovered based on various spectroscopic measurements, listed in Section 3."572" However, given the large uncertainties on the available data and the degeneracy between retrieved parameters, the existence of some of these species is called into question."," However, given the large uncertainties on the available data and the degeneracy between retrieved parameters, the existence of some of these species is called into question."573 We therefore set out to statistically test whether the presence of each molecule is truly warranted by the data., We therefore set out to statistically test whether the presence of each molecule is truly warranted by the data.574" Using a statistical significance test known as an F-test (Snedecor&Cochran1989), we evaluated the change in confidence level between a range of simple and complex models."," Using a statistical significance test known as an $F$ –test \citep{sne89}, we evaluated the change in confidence level between a range of simple and complex models."575 We successively added more complexity to the model(i.e., We successively added more complexity to the model(i.e.576 added molecules to the H2/He atmosphere to increase the number of degrees of freedom) and evaluated the improvement in the goodness-of-fit parameter (x?) for the best-fitting spectrum., added molecules to the $_{2}$ /He atmosphere to increase the number of degrees of freedom) and evaluated the improvement in the goodness-of-fit parameter $\chi^2$ ) for the best-fitting spectrum.577 The F-test was then used to assess whether the inclusion of a particular molecule was required to fit the data., The $F$ –test was then used to assess whether the inclusion of a particular molecule was required to fit the data.578 The F'-test indicates that an atmospheric composition for HD 189733b containing H2O and CO» are highly plausible at the 2.99.9846 confidence level., The $F$ –test indicates that an atmospheric composition for HD 189733b containing $_{2}$ O and $_{2}$ are highly plausible at the $>$ confidence level.579 The presence of these two molecules is required to obtain a reasonable fit to the available data., The presence of these two molecules is required to obtain a reasonable fit to the available data.580" In addition, we find insignificant confidence levels (i.e. <95%)) when we increase the complexity of the model by adding CO, CH, or a combination of the two to the simple model containing H», He, H30 and CO»."," In addition, we find insignificant confidence levels (i.e. $<$ ) when we increase the complexity of the model by adding CO, $_4$ or a combination of the two to the simple model containing $_{2}$, He, $_{2}$ O and $_{2}$."581" In other words, CO and CH, provide negligible enhancements to the fitting quality and are not required to fit the dayside emission measurements to within the stated error bars (although upper limits on these species can certainly be derived)."," In other words, CO and $_{4}$ provide negligible enhancements to the fitting quality and are not required to fit the dayside emission measurements to within the stated error bars (although upper limits on these species can certainly be derived)."582" Additionally, we consider the implications of Gibsonetal. (2011), who claimed larger uncertainties on the HST/NICMOS data by re-processing its transmission spectrum (Swainetal. 2008)."," Additionally, we consider the implications of \citet{gib11}, who claimed larger uncertainties on the $HST$ /NICMOS data by re-processing its transmission spectrum \citep{swa08}."583". If the same conclusions are applicable to the secondary eclipse emission spectrum, then we must similarly increase the measurement error on the NICMOS data by a factor of five."," If the same conclusions are applicable to the secondary eclipse emission spectrum, then we must similarly increase the measurement error on the NICMOS data by a factor of five."584" This leads to even smaller confidence levels on the more complex models, and the solutions are found to be even more degenerate."," This leads to even smaller confidence levels on the more complex models, and the solutions are found to be even more degenerate."585" In this case, evaluating the F-test significance using a variety of models still suggests that both H2O and CO» are required to reproduce the dayside emission spectra (799.2946 confidence), but makes the presence of CO and CH, even more uncertain."," In this case, evaluating the $F$ –test significance using a variety of models still suggests that both $_{2}$ O and $_{2}$ are required to reproduce the dayside emission spectra $>$ confidence), but makes the presence of CO and $_4$ even more uncertain."586" We therefore conclude that, irrespective of the uncertainties on the H ST/NICMOS data, current datasets are unable to provide detections of CO and CH, on the dayside of HD 189733b with any sort of statistical certainty."," We therefore conclude that, irrespective of the uncertainties on the $HST$ /NICMOS data, current datasets are unable to provide detections of CO and $_4$ on the dayside of HD 189733b with any sort of statistical certainty."587" Nevertheless, upper limits on the abundances of these molecules can be obtained (see Section 5), and all four molecules will be included in our subsequent study."," Nevertheless, upper limits on the abundances of these molecules can be obtained (see Section 5), and all four molecules will be included in our subsequent study."588" Using the NEMESIS algorithm, we retrieve the best-fitting dayside spectrum of HD 189733b, incorporating both the Spitzer and HST observations as stated in Section 3."," Using the NEMESIS algorithm, we retrieve the best-fitting dayside spectrum of HD 189733b, incorporating both the $Spitzer$ and $HST$ observations as stated in Section 3."589 Fig., Fig.590" 1 demonstrates the best-fitting spectrum to these measurements, and the retrieved atmospheric P-T profile and molecular abundances for H2O, CO», CO, and CHa."," \ref{f1} demonstrates the best-fitting spectrum to these measurements, and the retrieved atmospheric $P$ $T$ profile and molecular abundances for $_{2}$ O, $_{2}$ , CO, and $_{4}$ ."591 Each panel will be described below., Each panel will be described below.592 Fig., Fig.593 2 shows the contributions from the four main gases included in our model., \ref{f2} shows the contributions from the four main gases included in our model.594 This figure also shows the wavelength ranges where the molecular contributions are distributed by, This figure also shows the wavelength ranges where the molecular contributions are distributed by595"Observatory gave the first clues of a potential liuk between crystalline material in protoplanetary disks aud conmets,", gave the first clues of a potential link between crystalline material in protoplanetary disks and comets.596 A ercat similarity was noted between the spectra of the disk around the Ierbig star IID 100516 aud that of comet IIale-Dopp (Crovisieretal.1997:Malfüitct1998).," A great similarity was noted between the spectra of the disk around the Herbig star HD 100546 and that of comet Hale-Bopp \citep{CR97,MA98}."597. More recently. the InfraRed Spectrograph (IRS. 5 38 qun. THoucketal. 20013) on-board theTelescope allowed an wuprecedcuted combination of high sensitivity aud the ability to observe large uunibers of disks. down to the brown dwiurf limit.," More recently, the InfraRed Spectrograph (IRS, 5 – 38 $\mu$ m, \citealt{HO04}) ) on-board the allowed an unprecedented combination of high sensitivity and the ability to observe large numbers of disks, down to the brown dwarf limit."598 The shape of the silicate features probed by the IRS spectra at 10 and 20 san is affected by the composition. size and structure of its cimitting dust.," The shape of the silicate features probed by the IRS spectra at 10 and 20 $\mu$ m is affected by the composition, size and structure of its emitting dust."599 Amorplous silicates show broad smooth uid-IR features. while the opacities of crystalline graius show sharp features duc to their larec-scale lattice arrangcuient. such that even small fractious of crystalline erains produce additional structure in the silicate features (Ainctal.2005:Bouwiuanal.2008:πιάot2009:Olofssonet 2010).," Amorphous silicates show broad smooth mid-IR features, while the opacities of crystalline grains show sharp features due to their large-scale lattice arrangement, such that even small fractions of crystalline grains produce additional structure in the silicate features \citep{MI05,BO08,JU09,OF10}."600. Because inmost protoplauctary disks are optically thick at optical and IR wavelengths. the silicate features observed iu the mid-IR are generally emitted by dust in the optically thin disk surface oulv.," Because most protoplanetary disks are optically thick at optical and IR wavelengths, the silicate features observed in the mid-IR are generally emitted by dust in the optically thin disk surface only."601 To probe the disk iidplaue. observations at longer waveleueths are recessary.," To probe the disk midplane, observations at longer wavelengths are necessary."602 Additionally. the ciission at 10 and 20 jan as been shown to arise frou differeut erain populations. xobius differeut radii (IXessler-Silaccietal.2006:Olofssonetal.2009.Re] 2010)..," Additionally, the emission at 10 and 20 $\mu$ m has been shown to arise from different grain populations, probing different radii \citep{KE06,OF09,OF10}. ."603 While the 10 san feature probes a warluer dust population. at < 1 AU for T Tauri stars. he dust cinitting at 20 gan is colder. further out and deeper iuto the disk (vessler-Silaceietal.2007).," While the 10 $\mu$ m feature probes a warmer dust population, at $\leq$ 1 AU for T Tauri stars, the dust emitting at 20 $\mu$ m is colder, further out and deeper into the disk \citep{KE07}."604. Two incthods lave been proposed to explain the orlation of crystal erains: thermal annealing of amorphous erains or vaporization followed by gas-phase condensation., Two methods have been proposed to explain the formation of crystal grains: thermal annealing of amorphous grains or vaporization followed by gas-phase condensation.605 Both methods require high temperatures (above ~LOOO Is. Fabianetal.2000:ένα 20013) which is Inconsistent with outer disk temperatures.," Both methods require high temperatures (above $\sim$ 1000 K, \citealt{FA00,GA04}) ) which is inconsistent with outer disk temperatures."606 However. crystalline eraius have been observed iu outer. as well as dmn inner disks (vauBockeletal.2001).," However, crystalline grains have been observed in outer, as well as in inner disks \citep{VB04}."607. Large-scale radial mixing has been invoked to explain the presence of crystals at low temperatures iu the outer disk (Bockeléc-Alorviuetal.2000:Cail2001:Ciesla 2009)..," Large-scale radial mixing has been invoked to explain the presence of crystals at low temperatures in the outer disk \citep{BM00,GA04,CI09}. ."608 Α third proposed formation mechanisin for crystal formation is that shock waves could locally heat amorphous silicates and crystallize thei (Desch&Connolly2002:IkerDesch 2002).," A third proposed formation mechanism for crystal formation is that shock waves could locally heat amorphous silicates and crystallize them \citep{DC02,HD02}."609. Frou protoplanetary disks to comets. several autlors have attempted to infer the dust composition from IRS spectra aud laboratory data on amorphous and crystalline silicate dust. using a variety of analysis techniques.," From protoplanetary disks to comets, several authors have attempted to infer the dust composition from IRS spectra and laboratory data on amorphous and crystalline silicate dust, using a variety of analysis techniques."610 Whether for individual objects etal 2008).. for mixed disk suples (Bowmanctal. 2010).. or systematic studies of the disk population of a given star-forming region (Sicilia-Agnilayctal.2009:Watsonctal.2009:Sargeutet 2009).. it has heen shown that a significant mass fraction of the dust in those disks waist be iun crystalline form.," Whether for individual objects \citep{FO04,ME07,PI08,BY08}, for mixed disk samples \citep{BO01,AP05,VB05,BO08,OF09,OF10,JU10}, or systematic studies of the disk population of a given star-forming region \citep{SI09,WA09,ST09}, it has been shown that a significant mass fraction of the dust in those disks must be in crystalline form."611 However. the many studies dealing with the uuneralogical composition of dust to date focus ou a specific region or object. füling to investigate the hypothesis that the crystallinity fraction Is a 1ucasure of the evolutionary stage of a region.," However, the many studies dealing with the mineralogical composition of dust to date focus on a specific region or object, failing to investigate the hypothesis that the crystallinity fraction is a measure of the evolutionary stage of a region."612 That is. uo study in the literature has vet investigated au mcerease of crystallinity fraction with cluster age.," That is, no study in the literature has yet investigated an increase of crystallinity fraction with cluster age."613 Mineralogical studies of Solar Systei bodies show a range of crystallinity fractious., Mineralogical studies of Solar System bodies show a range of crystallinity fractions.614" Evidence frou, primitive chondrites shows that the abundance of crystalline silicate imaterial varies from nearly nothing wp to 20 S0 (em. Acter 091 aud ALTT7307. Poutoppidan&Brearley2010 and references therein)."," Evidence from primitive chondrites shows that the abundance of crystalline silicate material varies from nearly nothing up to 20 – 30 (e.g. Acfer 094 and ALH77307, \citealt{PB10}615 and references therein)."616 Oort cloud colucts. with long periods and large distances from the Sun. have interred crystallinity fractious up to 6 SO (e.@. TWale-Bopp. Woodenetal.1999.2007)). -Jupiter-zu.," Oort cloud comets, with long periods and large distances from the Sun, have inferred crystallinity fractions up to 60 – 80 (e.g. Hale-Bopp, \citealt{WO99,WO07}) ). Jupiter-family,"617 or short period comets. have lower fractions. up ο 735 (e.g. OP/Tempel 1. Warkeretal.2007:: 81P/Wild 2. Zoleuskyctal.2006)).," or short period comets, have lower fractions, up to $\sim$ 35 (e.g. 9P/Tempel 1, \citealt{HA07}; 81P/Wild 2, \citealt{ZO06}) )."618 This discrepancy dmi fractions ntfs to the existence of a radial dependence iu crystallinity fraction in the protoplanetary disk around he vouug Sun (Warkeretal.2005)., This discrepancy in fractions points to the existence of a radial dependence in crystallinity fraction in the protoplanetary disk around the young Sun \citep{HA05}.619. It is important to rote that those values are model dependent. aud the use of large amorphous erains (10 /— 100 jun) can lead ο systematically lower ervstalline. fractious (Ibarkeral.2002).," It is important to note that those values are model dependent, and the use of large amorphous grains (10 – 100 $\mu$ m) can lead to systematically lower crystalline fractions \citep{HA02}."620. This is evident for Hale-Dopp. where Miuotal.(2005). find a πιο lower fraction (7.5 %)) than other authors. using a distribution of amorphous erain sizes up to 100 yan. What is clear is that even within the discrepancies. the crystallinity fractions derived for Solar System bodies are appreciably higher than those derived for the ISM dust (<<2C. Isemperetal.2001).," This is evident for Hale-Bopp, where \citet{MI05a} find a much lower fraction $\sim$ 7.5 ) than other authors, using a distribution of amorphous grain sizes up to 100 $\mu$ m. What is clear is that even within the discrepancies, the crystallinity fractions derived for Solar System bodies are appreciably higher than those derived for the ISM dust $< 2 \%$, \citealt{KE04}) )."621 RecentSpitzer data indicate further similarities between crystalline silicate features seen dm comets or asteroids with those seen iu some debris disks around solar mass stars (Beichmanetal.2006:Lisse2007.2008).," Recent data indicate further similarities between crystalline silicate features seen in comets or asteroids with those seen in some debris disks around solar mass stars \citep{BE06,LI07,LI08}."622. One proposed explanation is that the observed spectral eatures m the disk result from the catastrophic break-up of a single huge bodv (a super comet) which creates the small dust particles needed. for detection., One proposed explanation is that the observed spectral features in the disk result from the catastrophic break-up of a single large body (a `super comet') which creates the small dust particles needed for detection.623 At the even emlier protoplauetarv disk stage. there is iuüted observational evidence for radial gradients iu crystallinity from mid-infrared interferometry data; with Heher crystallinity fractious found closer to the voune stars (vanBockelctal.2001:Schegereret2008).," At the even earlier protoplanetary disk stage, there is limited observational evidence for radial gradients in crystallinity from mid-infrared interferometry data, with higher crystallinity fractions found closer to the young stars \citep{VB04,SC08}."624. All of this suggests that the crystallization occurs carly in he disk evolution aud is then incorporated iuto larger solid bodies., All of this suggests that the crystallization occurs early in the disk evolution and is then incorporated into larger solid bodies.625 Besides dust composition. the evolution of grain sizes ds an essential indicator of disk evolution.," Besides dust composition, the evolution of grain sizes is an essential indicator of disk evolution."626 The initially sub-jun size ISAL grains unust grow astounding LL15 orders of mnaguitude in diuneter if they are o form planets., The initially $\mu$ m size ISM grains must grow astounding 14–15 orders of magnitude in diameter if they are to form planets.627 If grains were to grow orderly aud steadily. theoretical calculations predict disks to have ‘lly dissipated their small eraius within ~10° years (Weideuschilling1980:Diullemoud&Domini2005).," If grains were to grow orderly and steadily, theoretical calculations predict disks to have fully dissipated their small grains within $\sim$ $^5$ years \citep{WE80,DD05}."628. The fact that many disks a few Myv. old. are observed o have sinall erains (Hernandezetal.2008) poses a serious problem for the paradigni that eran erowth is a steady. monotonie process in disk evolution aud planotcorlmation.," The fact that many disks a few Myr old are observed to have small grains \citep{HE08} poses a serious problem for the paradigm that grain growth is a steady, monotonic process in disk evolution and planetformation."629 Additionally. s2all dust has been observed ii jesurface lavers of disks in clusters of different ages aud euvironnieuts for hundreds of svsteiis.," Additionally, small dust has been observed in thesurface layers of disks in clusters of different ages and environments for hundreds of systems."630 The implications. as discussed inost receutlv by Oliveiraetal.(2010) xl Olofssonetal. (2010)... is that sinall erains must be replenished by fragmentation of σος. grains. aud that 1equilibrium between eran growth aud fragmentation is established.," The implications, as discussed most recently by \citet{OL10} and \citet{OF10}, , is that small grains must be replenished by fragmentation of bigger grains, and that anequilibrium between grain growth and fragmentation is established."631 Oliveiraetal.(2010). have shown that this," \citet{OL10}632 have shown that this"633"was strongly modified (new opacities. new nuclear reaction rates, new equation of state. etc.).","was strongly modified (new opacities, new nuclear reaction rates, new equation of state, etc.)."634 The new model is different from that given in Bazot et al. (2005))., The new model is different from that given in Bazot et al. \cite{bazot05}) ).635 Apart from the constraint on the Y value. the basic reason is that in this previous paper the comparisons were only based on the stellar luminosity. which was misleading due to the Hipparcos parallax. which was later modified.," Apart from the constraint on the Y value, the basic reason is that in this previous paper the comparisons were only based on the stellar luminosity, which was misleading due to the Hipparcos parallax, which was later modified."636 Also the method used at that time was not as precise as the one we now use., Also the method used at that time was not as precise as the one we now use.637 Note that the scaling of parameters may lead to wrong results for stars in which the seismic modes cannot be precisely identified., Note that the scaling of parameters may lead to wrong results for stars in which the seismic modes cannot be precisely identified.638 This is the case for example for the mass proposed by Kallinger et al. RO , This is the case for example for the mass proposed by Kallinger et al. \cite{kallinger08}) )639for the star ze Arae. which is much too large (1.23 ," for the star $\mu$ Arae, which is much too large (1.23 )."640We also ο. an analysis of the size of the mixed core by testing the implications of overshooting on the mode frequencies., We also performed an analysis of the size of the mixed core by testing the implications of overshooting on the mode frequencies.641 We found a strong constraint on the possibility of core overshooting. treated as an extension of convection: the size of this extension must be less than 0.5 of the pressure scale height (overshooting parameter).," We found a strong constraint on the possibility of core overshooting, treated as an extension of convection: the size of this extension must be less than 0.5 of the pressure scale height (overshooting parameter)."642 This does not exclude other kinds of mild boundary effects at the edge of the core. provided that they do not lead to strong mixing.," This does not exclude other kinds of mild boundary effects at the edge of the core, provided that they do not lead to strong mixing."643 At the present time. we were able to perform this deep seismic analysis on two solar type stars hosting planets. both with a large metallicity (about twice solar): t Hor (Vauclair et al 2008)) and µ Arae (this paper).," At the present time, we were able to perform this deep seismic analysis on two solar type stars hosting planets, both with a large metallicity (about twice solar): $\iota$ Hor (Vauclair et al \cite{vauclair08}) ) and $\mu$ Arae (this paper)."644 In both cases. precise stellar parameters could be obtained.," In both cases, precise stellar parameters could be obtained."645 We found however an important difference between these two overmetallie stars., We found however an important difference between these two overmetallic stars.646 In ¢ Hor. the helium abundance ts low. even lower than the solar value. in accordance with the helium value determined for the Hyades stellar cluster.," In $\iota$ Hor, the helium abundance is low, even lower than the solar value, in accordance with the helium value determined for the Hyades stellar cluster."647 As other observational parameters also coincide. we concluded that¢ Hor is an ejected member of the Hyades.," As other observational parameters also coincide, we concluded that $\iota$ Hor is an ejected member of the Hyades."648 The reason why the helium abundance is so low in these stars while the metallicity is high ts still a mystery. although it certainly depends on the mass of the stars that polluted the original nebula.," The reason why the helium abundance is so low in these stars while the metallicity is high is still a mystery, although it certainly depends on the mass of the stars that polluted the original nebula."649 In i Arae. on the other hand. the helium abundance is large. as expected from the usual laws for the chemical evolution of galaxies (Isotov Thuan 2004)).," In $\mu$ Arae, on the other hand, the helium abundance is large, as expected from the usual laws for the chemical evolution of galaxies (Isotov Thuan \cite{isotov04}) )."650 This star was formed in a nebula which suffered normal pollution from stars with proportional yields of helium and metals., This star was formed in a nebula which suffered normal pollution from stars with proportional yields of helium and metals.651 Seismology can lead to precise values of helium abundances in solar type stars. where helium cannot be directly derived from spectroscopy.," Seismology can lead to precise values of helium abundances in solar type stars, where helium cannot be directly derived from spectroscopy."652 This represents a success. quite apart from all other results and constraints. and will be of importance for the study of the chemical evolution of our Galaxy.," This represents a success, quite apart from all other results and constraints, and will be of importance for the study of the chemical evolution of our Galaxy."653"he 2676 and 251 clusters using e,.=1000kms.+ for the galactic winds CTH-SW runs).",the g676 and g51 clusters using $v_w=1000 \vel$ for the galactic winds (TH-SW runs).654 However. while the effect of the stronger eedback is that of decreasing the number of galaxies. (see also Sect.," However, while the effect of the stronger feedback is that of decreasing the number of galaxies, (see also Sect."655 4.4. below). it leaves their metal content. and. therefore. the ugh CMR normalization. almost unchanged.," \ref{s:lf} below), it leaves their metal content, and, therefore, the high CMR normalization, almost unchanged."656 Although a Salpeter IMF fares rather well as for the CMR. we note that all the BCGs (big filled circles in Fig.2)) are much bluer. by about 0.5 magnitude. than expected from the red sequence.," Although a Salpeter IMF fares rather well as for the CMR, we note that all the BCGs (big filled circles in \ref{fi:cmr_z0}) ) are much bluer, by about 0.5 magnitude, than expected from the red sequence."657 Such a blue excess of the colors of the BGCs. which takes place despite heir high metallicity. finds its origin in the large star formation rate. associated to overeooling. which takes place in the central cluster regions.," Such a blue excess of the colors of the BGCs, which takes place despite their high metallicity, finds its origin in the large star formation rate, associated to overcooling, which takes place in the central cluster regions."658" Typical values for the star formation rate of the BCG in our simulations are in range 600—1000 3. /vr for the most massive clusters Mou,1077AIAE. rand ~1004. /yr for the least massive ones (Alou)c1071ΑΗ. 1."," Typical values for the star formation rate of the BCG in our simulations are in range 600–1000 $M_\odot$ /yr for the most massive clusters $M_{200}\simeq 10^{15}\msun$ ) and $\sim659100 M_\odot$ /yr for the least massive ones $M_{200}\simeq66010^{14}\msun$ )."661 Although observations indicate the presence of some ongoing star formation in some BCGs located at the center of cool core clusters. they are always at a much lower level and consistent with a star formation rate of  10-1004. /yr for clusters of comparable richness andreferences therein).," Although observations indicate the presence of some ongoing star formation in some BCGs located at the center of cool core clusters, they are always at a much lower level and consistent with a star formation rate of $\sim 10$ $100 M_\odot$ /yr for clusters of comparable richness \citep[e.g., ][and references662therein]{1987MNRAS.224...75J,2006astro.ph..2323B,2006astro.ph..4044M}."663 The effect of recent star formation on the CMR is explicitely shown in Figure 3.., The effect of recent star formation on the CMR is explicitely shown in Figure \ref{fi:cmr_z1}.664 We show here the case in which all star particles. formed at redshift 2 <1 are excluded from the computation of the galaxy luminosities.," We show here the case in which all star particles, formed at redshift $z<1$ are excluded from the computation of the galaxy luminosities."665 This is equivalent to assume that we completely quench star formation since 2=1., This is equivalent to assume that we completely quench star formation since $z=1$.666 Neglecting recent star formation has the twofold effect of reducing the scatter in the CMR and of making BCG colors signiticantly redder. although they still fall slightly below the observed relation.," Neglecting recent star formation has the twofold effect of reducing the scatter in the CMR and of making BCG colors significantly redder, although they still fall slightly below the observed relation."667 A numberobservational analyses have established that the mass—to-light ratioof in clusters generally increases with the cluster mass. AM/LxM with 2 0.2-0.4. over a fairly large dynamic range. from poor groups to rich clusters (e.g.2222227222).," A number of observational analyses have established that the mass--to--light ratio in clusters generally increases with the cluster mass, $M/L\propto M^\gamma$ with $\gamma \simeq 0.2$ –0.4, over a fairly large dynamic range, from poor groups to rich clusters \citep[e.g.,668][]{1998A&A...331..493A,2000ApJ...530...62G,2002ApJ...569..720G,2002ApJ...565L...5B,2003ApJ...591..749L,2004ApJ...610..745L,2004AJ....128.1078R,2004AJ....128.2022R,2005A&A...433..431P}."669 A likely explanation for this trend is the reduced cooling efficiency within more massive. hotter halos (e.g...2)... which reduces star formation within richer clusters.," A likely explanation for this trend is the reduced cooling efficiency within more massive, hotter halos \citep[e.g.,670][]{2003MNRAS.339..312S}, which reduces star formation within richer clusters."671 In fact. an increasing trend of AML with cluster mass is naturally predicted by semi-analytical models of galaxy formation (e.g...2)..," In fact, an increasing trend of $M/L$ with cluster mass is naturally predicted by semi–analytical models of galaxy formation \citep[e.g., ][]{1999MNRAS.303..188K}."672 In Figure 4. we compare the relation between mass and luminosity within rou for our simulated clusters. and compare it to the ;-band results by ?. and to the A—band results by ?..," In Figure \ref{fi:ml} we compare the relation between mass and luminosity within $r_{500}$ for our simulated clusters, and compare it to the $i$ –band results by \cite{2005A&A...433..431P} and to the $K$ –band results by \cite{2004ApJ...610..745L}."673 In general. we find that the A//2 from simulations is rather close to the observed one in the / band. also with a comparably small scatter.," In general, we find that the $M/L$ from simulations is rather close to the observed one in the $i$ band, also with a comparably small scatter."674 In the A band. a Salpeter IMF still agrees with observations within the statistical uncertainties. while the top-heavy IMFproduces too red galaxies. thus consistent with the results of the CMR. as shown in Fig.2..," In the $K$ band, a Salpeter IMF still agrees with observations within the statistical uncertainties, while the top–heavy IMFproduces too red galaxies, thus consistent with the results of the CMR, as shown in \ref{fi:cmr_z0}."675" We fit our mass—luminosity relation with a power-law——————)'.. we find (0.3);=(0.74.0.92) and (a.3),=(0.76.3.2) in the 7 and ἐν band. respectively. for the runs with Salpeter IMF. while (a.:7);=(0.70.0.91) and (a),=(0.74.4.7) or the top-heavy IMF."," We fit our mass–luminosity relation with a power–law, we find $(\alpha,\beta)_i=(0.74,0.92)$ and $(\alpha,\beta)_K=(0.76,3.2)$ in the $i$ and $K$ band, respectively, for the runs with Salpeter IMF, while $(\alpha,\beta)_i=(0.70,0.91)$ and $(\alpha,\beta)_K=(0.74,4.7)$ for the top–heavy IMF."676 Therefore. our simulations agree with qe observational trend for an increasing mass-to-light ratio with cluster mass. independent of the IMF and luminosity band.," Therefore, our simulations agree with the observational trend for an increasing mass-to-light ratio with cluster mass, independent of the IMF and luminosity band."677 Using qe stronger feedback for the top-heavy IMF turns into a sizeable ραuppression of the luminosity. especially for 51.," Using the stronger feedback for the top–heavy IMF turns into a sizeable suppression of the luminosity, especially for g51."678 The reasonable level of agreement between the observed and je simulated AZ/£L may suggest that our simulations produces a realistic population of galaxies., The reasonable level of agreement between the observed and the simulated $M/L$ may suggest that our simulations produces a realistic population of galaxies.679 However. as demonstrated in Figure 5.. this is not the case.," However, as demonstrated in Figure \ref{fi:ngal}, , this is not the case."680 In this figure. we compare the simulated and observed number of cluster galaxies. brighter than a given luminosity limit. both in ¢ and in the A bands.," In this figure, we compare the simulated and observed number of cluster galaxies, brighter than a given luminosity limit, both in $i$ and in the $K$ bands."681 Clearly. simulations underpredict such a number. by a factor ~ 2—3.," Clearly, simulations underpredict such a number, by a factor $\sim 2$ –3."682 This result is at variance with respect to that from semi-analytical models of galaxy formation. which instead predict the correct number of cluster members (e.g.. 22Y.. ," This result is at variance with respect to that from semi–analytical models of galaxy formation, which instead predict the correct number of cluster members \citep[e.g.,][]{2004MNRAS.349.1101D,2005MNRAS.361..369L}. ."683However. semi-analytical models are generally successfulin producing the correct LF.," However, semi–analytical models are generally successfulin producing the correct LF."684 They employ a suitable technique to track galaxies. based on the," They employ a suitable technique to track galaxies, based on the"685We have analysed the properties of pairs of galaxies in high ensitv environments corresponding to eroups and. clusters X ealaxies with virial masses 1077105734...,We have analysed the properties of pairs of galaxies in high density environments corresponding to groups and clusters of galaxies with virial masses $10^{13} - 10 ^{15} M_{\odot}$.686 We stress the unact that the analysis discussed in this paper is based on —16 comparison between two galaxy. samples which cdilfered between cach otherο on the [act that one of sample 'omprises galaxies with close companions. and the other. o not.," We stress the fact that the analysis discussed in this paper is based on the comparison between two galaxy samples which differed between each other on the fact that one of sample comprises galaxies with close companions, and the other, do not."687 llence. we are always estimating the statistical ilferences in the properties of galaxies introduced by the presence of a companion.," Hence, we are always estimating the statistical differences in the properties of galaxies introduced by the presence of a companion."688 The main conclusions can be summarized as follow:, The main conclusions can be summarized as follow:689"the dust continuum emission is detected at ""mm and al 10n (IIughes et 22007: Ratzka et 22007).",the dust continuum emission is detected at 7mm and at $10\micron$ (Hughes et 2007; Ratzka et 2007).690" such a dust distribution could be interpreted as a result of grainMN (Strom et 11939: Dullemond Dominik 2005). (Clarke et-- ""22001: Alexander et 22006: Gorti ILollenbach 2009: OwenphotevaporatIon et 22010). or a gap an orbiting giant planet (e.g.. Marsh Mahoney. 1992: Calvet et 22002: Rice et 22003: Quillen et 2004: Calvet et 22005)."," Such a dust distribution could be interpreted as a result of grain growth (Strom et 1989; Dullemond Dominik 2005), photoevaporation (Clarke et 2001; Alexander et 2006; Gorti Hollenbach 2009; Owen et 2010), or a gap opened by an orbiting giant planet (e.g., Marsh Mahoney 1992; Calvet et 2002; Rice et 2003; Quillen et 2004; Calvet et 2005)."691 These scenarios could potentially be distinguished by examining properties other than the SED. since the scenarios make different predictions for e stellar accretion rate. disk mass. and radial distribution of gas in the disk (Najita et 22007a: Alexander Armitage 2007).," These scenarios could potentially be distinguished by examining properties other than the SED, since the scenarios make different predictions for the stellar accretion rate, disk mass, and radial distribution of gas in the disk (Najita et 2007a; Alexander Armitage 2007)."692" The first two diagnostics (stellar aceretion rates and disk masses) were used by Najita et ((2007a)"" to probe (he nature of transition objects in the Taurus star forming region.", The first two diagnostics (stellar accretion rates and disk masses) were used by Najita et (2007a) to probe the nature of transition objects in the Taurus star forming region.693 They found Taurus transition objects have higher than average disk masses as well as stellar accretion rates that are ~10 times lower (han non-transition objects., They found that Taurus transition objects have higher than average disk masses as well as stellar accretion rates that are $\sim 10$ times lower than non-transition objects.694 These properties are roughly consistent with the predictions of theories of giant planet formation (e.g.. Lubow οἱ 11999: Lubow D'Angelo 2006).," These properties are roughly consistent with the predictions of theories of giant planet formation (e.g., Lubow et 1999; Lubow D'Angelo 2006)."695" The high disk mass of TW Ilva (>0.06.: Calvet el 22002) and its comparatively low stellar accretion rate (~10.9M,vr.ts e.g. Herczeg οἱ 22004: Muzerolle et 22000: Alenear Dasri 2000) place it in a similar region of the M Maik plane as (he Taurus transition objects."," The high disk mass of TW Hya $> 0.06\Msun$; Calvet et 2002) and its comparatively low stellar accretion rate $\sim 10^{-9}\Msunperyr$; e.g., Herczeg et 2004; Muzerolle et 2000; Alencar Basri 2000) place it in a similar region of the $\Mdot$ $\Mdisk$ plane as the Taurus transition objects."696 sS(udies of line emission Irom the gaseous component of (rausition disks. like that presented. here. offer (he opportunity (to complement studies of stellar accretion rates ancl disk mass. bv probing Che radialdistribution of gas in the disk and therefore the evolutionary state of the svstem.," Studies of line emission from the gaseous component of transition disks, like that presented here, offer the opportunity to complement studies of stellar accretion rates and disk mass, by probing the radial distribution of gas in the disk and therefore the evolutionary state of the system."697" As described byNajita et ((2007a. 2008): (1) In the grain growth: ancl planetesimal formation scenario. (he inner disk is rendered oplically thin in the continuum. but the gaseous component is unaltered and would fill the region within 25,4."," As described by Najita et (2007a, 2008): (1) In the grain growth and planetesimal formation scenario, the inner disk is rendered optically thin in the continuum, but the gaseous component is unaltered and would fill the region within $R_{\rm hole}$."698 Emission Irom gas witliin yo. would produce bright emission because of the lack of continuum emission [rom same region of the disk. (, Emission from gas within $R_{\rm hole}$ would produce bright emission because of the lack of continuum emission from same region of the disk. (6992) La planet has formed with a mass sullicient to open a gap (~ 11). gas will be cleared in the vieinity of its orbit. bul gap-crossing streams. from (he outer disk to the planet. and from the planet to the inner disk. can allow continued accretion onto both the planet and the star. the latter via (he replenishimment of theinner disk n uae€Bye ie.g.. Lubow et 11999: Ixlev 1999: Bryden et 11999: D'Angelo et no22003: shBate et ni22003: Lubow D'Angelo 2006).,"2) If a planet has formed with a mass sufficient to open a gap $\sim 1 M_J$ ), gas will be cleared in the vicinity of its orbit, but gap-crossing streams, from the outer disk to the planet, and from the planet to the inner disk, can allow continued accretion onto both the planet and the star, the latter via the replenishment of the inner disk within $R_{\rm inner} < R_{\rm hole}$ (e.g., Lubow et 1999; Kley 1999; Bryden et 1999; D'Angelo et 2003; Bate et 2003; Lubow D'Angelo 2006)."700 While the ihen guüseotts enission lines. the low surlace filling [actor of gas in the region of the gap would produce weak to negligible emission because of the small. projected emitting area of the accretion," While the might then produce gaseous emission lines, the low surface filling factor of gas in the region of the gap would produce weak to negligible emission because of the small projected emitting area of the accretion"701(53. atl o describe the planets and planetary svstenis emerging roii the| Eta-EarthEta-Ear Progra.,$\sim$ $\sim$ to describe the planets and planetary systems emerging from the Eta-Earth Program.702OCT: Thle statisticstatith off plananetf occi:occur-o hfroii thee EtaEta-EarEarth Survey will offer⋅ important constraints ou compctine heories∖⊳⋠∖⊳⋅∢⋅ of planct formation (77777)...2999)," The statistics of planet occurance from the Eta-Earth Survey will offer important constraints on competing theories of planet formation \citep{Ida04a,Kenyon06,Alibert05,Mordasini07,Ida_Lin08_iv}."703" Thi↴* modηςIs differ in assuniptious about the erowth rate of dust iuto ο”.B we DUB otB n i. M . anc t o the snow Bue qm the ME the cflicacy o te ectsDea the MEM ot ent aud “Ter. aud te ""nrelevance of plauet-planct mteractions."," These models differ in assumptions about the growth rate of dust into planetesimals, the viscosity of the disk, the location and effects of the snow line in the disk, the efficacy of inward migration, the accretion of gas and water, and the relevance of planet-planet interactions."704 lY Gut consistent with detections and measurements of joviau gas glauuts (og.TT).," They are consistent with detections and measurements of jovian gas giants \citep[e.g.][]{Marcy_Japan_05,Udry2003}."705" However. these theories predict hat planets of mass are rare within AAU. forming a ""planet desert? of super-Earths (Mj sinis nearthe)) aud Neptune-mass planets."," However, these theories predict that planets of mass are rare within AU, forming a “planet desert” of super-Earths $M_{\mathrm{pl}}$ $i$ $\le$ ) and Neptune-mass planets."706" All models predict hat Type I inigration quickly (7 ~ 110 vr) causes the rocky planets to spiral iuxvard. destined to be lost in hestar""."," All models predict that Type I migration quickly $\tau$ $\sim$ $^5$ yr) causes the rocky planets to spiral inward, destined to be lost in the."707 Moeamwhile. the more massive rocky cores (AL LII incarthe)yucerctegasquicklg.becomingict Or(gus giantsD," Meanwhile, the more massive rocky cores $M$ $>$ ) accrete gas quickly, becoming ice- or gas-giants."708E heresultingdistributionof planctargsystemsisdepletedof pl iicarthwithinl AAG, The resulting distribution of planetary systems is depleted of planets in the mass range within AU.709", If the prediction of a low-niass desert is contradicted wea statistically: well-definedd sample oft stars (such as the Eta-Earth Survey).⊲⋅ then planet formation. theory must )o Siguificauthy modified with new phvses."," If the prediction of a low-mass desert is contradicted by a statistically well-defined sample of stars (such as the Eta-Earth Survey), then planet formation theory must be significantly modified with new physics."710 The planet desert has indeed∙ been challenged. bv?E who estimate: hataM ο ofD GI&paTp- dwiufs. have rocky or Noeptuneauass- anets miwud of 50dd orbits., The planet desert has indeed been challenged by \citet{Mayor09} who estimate that $\pm$ of GK dwarfs have rocky or Neptune-mass planets inward of d orbits.711 Thisclaim is based in wart on detections. bv the Swiss». Peroup. of. three Neptune--sizeH planets orbitingtat 660830⋅⋅∖∙ Li (23.. two oft three plaucts," Thisclaim is based in part on detections by the Swiss group of three Neptune-size planets orbiting 69830 \citep{Lovis2006}, , two of three planets"712suggests strongly for high-order atmospheric gravity waves and they can be excited by turbulent stresses in the convection zone.,suggests strongly for high-order atmospheric gravity waves and they can be excited by turbulent stresses in the convection zone.713 We can summarize the main results derived from the analysis of 1600 A continuum observations are as follows: (1) The uv bright points. uv network elements and uv background regions will exhibit a fluctuations with a smaller period in their intensity oscillations. (," We can summarize the main results derived from the analysis of 1600 $\AA$ continuum observations are as follows: (i) The uv bright points, uv network elements and uv background regions will exhibit a fluctuations with a smaller period in their intensity oscillations. ("71411) We find evidence from the power spectrum analysis for a longer period of oscillations: the uv bright points are associated with around 5.5 hours. the uv network elements exhibit around 4.6 hours and whereas the background regions show around 3.4 hours. (,"ii) We find evidence from the power spectrum analysis for a longer period of oscillations: the uv bright points are associated with around 5.5 hours, the uv network elements exhibit around 4.6 hours and whereas the background regions show around 3.4 hours. ("71511) It is noted that the different,iii) It is noted that the different716Eclipsiug Binaries (hereafter EBs) are excelleut objects for determining the physical properties of stars and detecting additional components iu them.,Eclipsing Binaries (hereafter EBs) are excellent objects for determining the physical properties of stars and detecting additional components in them.717 The loug-time behavior of the period of au ED could reveal the presence of another component orbiting with the ED arouud the conuimion center of nass;, The long-time behavior of the period of an EB could reveal the presence of another component orbiting with the EB around the common center of mass.718 Photometric observations of EBs soluctimes cover more than a ceutury. therefore it is possible to detect the third bodies with a similar period.," Photometric observations of EBs sometimes cover more than a century, therefore it is possible to detect the third bodies with a similar period."719 The motion around the barvcenter causes apparent changes of the observed. binary’s period with a period corresponding to the orbital one of the third body. called he LiIeht-Time Effect (or light-travel time’. hereafter LITE).," The motion around the barycenter causes apparent changes of the observed binary's period with a period corresponding to the orbital one of the third body, called the LIght-Time Effect (or 'light-travel time', hereafter LITE)."720 Lewin(1959). nuproved the method developed w Woltjer(1922) for analyzing the long-term variation of the times of ruininia caused by a third body orbiting he eclipsing pair., \cite{Irwin1959} improved the method developed by \cite{Woltjer1922} for analyzing the long-term variation of the times of minima caused by a third body orbiting the eclipsing pair.721 Useful conmucuts and limitations were discussed. bv Fricboes-Coude&Terczee(1973) aud w Maver(010990)., Useful comments and limitations were discussed by \cite{FCH73} and by \cite{Mayer1990}.722.. Nowadays there are more than one mudred EBs showing LITE. where the effect is certainly preseuted or supposed (see e.g. Borkovits&Ilegedüs (1996).. Albavraketal. (1999).. Wolfetal. (2001).. Ποιαetal. (2006).. 0te.).," Nowadays there are more than one hundred EBs showing LITE, where the effect is certainly presented or supposed (see e.g. \cite{BorkovitsHegedus}, \cite{Albayrak1999}, \cite{Wolf2004}, , \cite{Hoffman2006}, , etc.)."723 See the catalogue of the οC diagrauis by Kreineretal.(2001).. where the apparent orbital period. changes im many EBs are preseuted.," See the catalogue of the $O-C$ diagrams by \cite{Kreiner2001}, where the apparent orbital period changes in many EBs are presented."724 The look of OC diagrams in the present study was adopted to be the same as in this catalogue., The look of $O-C$ diagrams in the present study was adopted to be the same as in this catalogue.725 In our figures to 8 the full circles represent the primary and the open 1.circles the secondary times of minima. the bigger the point. the bigecr the weight.," In our figures \ref{FigWZAnd1} to \ref{FigAHTau} the full circles represent the primary and the open circles the secondary times of minima, the bigger the point, the bigger the weight."726 For the Iunitatious auc consequences of the O6 diagram analysis. see e.g. Sterken(2005).," For the limitations and consequences of the $O-C$ diagram analysis, see e.g. \cite{Sterken2005}."727" The computation of the paramcters of the third-body orbit is a classical inverse problem with 5 parameters to be fouud pe du ow. 63. which indicate the period of the third body. the periastrou passage. the seimu-auuplitude of the light-time effect. the argunoeut of periastron and the ουσΙΤ, respectively (for a detailed description see e.g. Mawer 19903)."," The computation of the parameters of the third-body orbit is a classical inverse problem with 5 parameters to be found – $p_3$, $T_0$, $A$, $\omega$, $e_3$, which indicate the period of the third body, the periastron passage, the semi-amplitude of the light-time effect, the argument of periastron and the eccentricity, respectively (for a detailed description see e.g. \citealt{Mayer1990}) )."728" The ephemerides for the individual svstems (D, aud P for the linear one and 4 for the quadratic ono) have to be calculated together with the parnueters of LITE.", The ephemerides for the individual systems $JD_0$ and $P$ for the linear one and $q$ for the quadratic one) have to be calculated together with the parameters of LITE.729" The amass function Ελ} and the minimal massof the thirdcomponent ALji,=Ms;-:snu/sy (for"," The mass function $f(M_3)$ and the minimal massof the thirdcomponent $M_{3,min} = M_3 \cdot \sin i_3$ (for"730for the formation of a detonation.,for the formation of a detonation.731 As the expected helium layer is located exactly in the region where the merger ts most violent and the hot-spots form. it may facilitate a detonation at even lower densities in our scenario.," As the expected helium layer is located exactly in the region where the merger is most violent and the hot-spots form, it may facilitate a detonation at even lower densities in our scenario."732 In summary. there are reasons to believe that detonations may form for even a wider range of conditions than the ones we have assumed here.," In summary, there are reasons to believe that detonations may form for even a wider range of conditions than the ones we have assumed here."733 The violent merger scenario can only be realized frequently enough to account for a significant fraction of all type la supernovae if it also works for moderate mass differences between the two white dwarfs., The violent merger scenario can only be realized frequently enough to account for a significant fraction of all type Ia supernovae if it also works for moderate mass differences between the two white dwarfs.734 Thus. we have to check how changing the mass ratio for binary systems with a fixed primary white dwarf mass affects the conditions during the merger.," Thus, we have to check how changing the mass ratio for binary systems with a fixed primary white dwarf mass affects the conditions during the merger."735 Only if mergers with a reasonable range also lead to the formation of a detonation. the scenario will be more than an exotic possibility.," Only if mergers with a reasonable range also lead to the formation of a detonation, the scenario will be more than an exotic possibility."736 To test this. we compare four different mergers.," To test this, we compare four different mergers."737" All have the same primary mass of 0.9Μ... but different secondary masses of 0.89 M... 0.81M... 0.76Mo. and 0.7Μα which give mass ratios of 0.99, 0.9. 0.84 and 0.78."," All have the same primary mass of $0.9\, \mathrm{M_\odot}$, but different secondary masses of $0.89\, \mathrm{M_\odot}$ , $0.81\, \mathrm{M_\odot}$, $0.76\, \mathrm{M_\odot}$ and $0.7\,738 \mathrm{M_\odot}$ which give mass ratios of $0.99$, $0.9$, $0.84$ and $0.78$."739 In our simulations. they are set up with initial periods of 258. 33s. 305. and 40s. respectively.," In our simulations, they are set up with initial periods of $25\,\mathrm{s}$, $33\,\mathrm{s}$, $36\,\mathrm{s}$, and $40\,\mathrm{s}$, respectively."740 Figure 4. shows densities and temperatures of all particles of these simulations at the time when the conditions are most favorable for a detonation as well as temperature slices through the centers of the binaries., Figure \ref{fig:qtest09} shows densities and temperatures of all particles of these simulations at the time when the conditions are most favorable for a detonation as well as temperature slices through the centers of the binaries.741 Obviously. there are considerable differences between these three systems.," Obviously, there are considerable differences between these three systems."742 Firstly. with decreasing mass ratio. the merger becomes less violent.," Firstly, with decreasing mass ratio, the merger becomes less violent."743" While the mergers with mass ratios of 0.99, 0.9 and 0.84 produce several hot particles that ignite carbon and reach temperatures above 2-K.IO? this is not the case for the merger with the smallest mass ratio."," While the mergers with mass ratios of $0.99$, $0.9$ and $0.84$ produce several hot particles that ignite carbon and reach temperatures above $2 \cdot 10^9744 \mathrm{K}$, this is not the case for the merger with the smallest mass ratio."745 It is still possible that resolving the interaction region better will also show hotter particles., It is still possible that resolving the interaction region better will also show hotter particles.746 However. at the moment it seems more likely that below a certain mass ratio the merger is just not violent enough to ignite a detonation.," However, at the moment it seems more likely that below a certain mass ratio the merger is just not violent enough to ignite a detonation."747 This suggests a limiting mass ratio for the violent merger scenario of around 0.8., This suggests a limiting mass ratio for the violent merger scenario of around $0.8$.748 Another difference is the dynamical effect of the merger on the primary white dwarf., Another difference is the dynamical effect of the merger on the primary white dwarf.749 In a nearly equal mass merger the primary star is heavily distorted. very similar to the equal mass merger described in ?..," In a nearly equal mass merger the primary star is heavily distorted, very similar to the equal mass merger described in \citet{pakmor2010a}."750 For a mass ratio of 0.9. the primary white dwarf remains unaffected in the center. but its surface Is distorted.," For a mass ratio of $0.9$, the primary white dwarf remains unaffected in the center, but its surface is distorted."751 For the smaller mass ratios. 1t stays completely intact and cool and is surrounded by the material of the disrupted less massive companion.," For the smaller mass ratios, it stays completely intact and cool and is surrounded by the material of the disrupted less massive companion."752 The etfect of the merger on the primary white dwarf can also be seer in the right panels of Figure 4.., The effect of the merger on the primary white dwarf can also be seen in the right panels of Figure \ref{fig:qtest09}.753 In the q=0.99 merger the central density of the remaining white dwarf is lower than in the other cases.," In the $\mathrm{q} =754 0.99$ merger the central density of the remaining white dwarf is lower than in the other cases."755 As for these mergers the amount of *°Ni produced depends sensitively on the central density of the remaining white dwarf. this leads to smaller °°Ni masses in mergers with higher mass ratios.," As for these mergers the amount of $^{56}\mathrm{Ni}$ produced depends sensitively on the central density of the remaining white dwarf, this leads to smaller $^{56}\mathrm{Ni}$ masses in mergers with higher mass ratios."756 While this does not change the scenario fundamentally. it breaks the relation between mass of the primary white dwarf and the final °°Ni mass of the explosion for dim explosions.," While this does not change the scenario fundamentally, it breaks the relation between mass of the primary white dwarf and the final $^{56}\mathrm{Ni}$ mass of the explosion for dim explosions."757" For explosions of mergers of more massive white dwarfs. however. where most of the ""Ni is produced in nuclear statistical equilibrium. this will be a minor effect only."," For explosions of mergers of more massive white dwarfs, however, where most of the $^{56}\mathrm{Ni}$ is produced in nuclear statistical equilibrium, this will be a minor effect only."758 As shown in Figure 4.. with decreasing mass ratio the merger becomes less violent and the density m the hot spots drops.," As shown in Figure \ref{fig:qtest09}, with decreasing mass ratio the merger becomes less violent and the density in the hot spots drops."759 This may make the formation of a detonation more dificult for smaller mass ratios., This may make the formation of a detonation more difficult for smaller mass ratios.760 For the smallest mass ratio no violent merger occurs and we see the onset of a different merger regime in which an accretion disk forms around primary white dwarf., For the smallest mass ratio no violent merger occurs and we see the onset of a different merger regime in which an accretion disk forms around primary white dwarf.761 Figure 5. shows the ejecta composition of the thermonuclear explosion of the merger of two 0.859M.. as described in ?.. which is close to the merger with a mass ratio of 0.99 described above.," Figure \ref{fig:comp} shows the ejecta composition of the thermonuclear explosion of the merger of two $0.89\, \mathrm{M_\odot}$ as described in \citet{pakmor2010a}, which is close to the merger with a mass ratio of $0.99$ described above."762 The final composition contains 0.03M.« of carbon. 0.54M. of oxygen. 1.05M. of intermediate-mass elements and 0.1M. of iron group elements.," The final composition contains $0.03 \, \mathrm{M_\odot}$ of carbon, $0.54 \, \mathrm{M_\odot}$ of oxygen, $1.05 \, \mathrm{M_\odot}$ of intermediate-mass elements and $0.1 \, \mathrm{M_\odot}$ of iron group elements."763 As the densities at the time the explosion happens do not exceed 2x10’gem™. nuclear burning does not reach uclear statistical equilibrium anywhere in the object.," As the densities at the time the explosion happens do not exceed $2 \times 10^7 \mathrm{g\ cm^{-3}}$, nuclear burning does not reach nuclear statistical equilibrium anywhere in the object."764 Instead. iron group elements are only produced by incomplete silicon burning.," Instead, iron group elements are only produced by incomplete silicon burning."765 As the nuclear burning in the detonation takes place at low densities. electron captures are not efficient.," As the nuclear burning in the detonation takes place at low densities, electron captures are not efficient."766" Therefore the initial electron fraction Y, of the unburned material 15 conserved throughout the nuclear burning.", Therefore the initial electron fraction $Y_e$ of the unburned material is conserved throughout the nuclear burning.767" In case of Y,=0.5 the iron group elements consist of almost pure ??Ni.", In case of $Y_e = 0.5$ the iron group elements consist of almost pure $^{56}\mathrm{Ni}$.768 A lower initial Y; is equivalent to more neutron-rich material in the pre-explosion composition (i.e. more Ne resulting from the He burning phase)., A lower initial $Y_e$ is equivalent to more neutron-rich material in the pre-explosion composition (i.e. more $^{22}$ Ne resulting from the He burning phase).769 It leads to the production of some amount of stable iron replacing some of the ~°Ni., It leads to the production of some amount of stable iron replacing some of the $^{56}\mathrm{Ni}$.770 In total. however. the fraction of stable iron group elements produced is very small.," In total, however, the fraction of stable iron group elements produced is very small."771 As shown in Figure 5.. iron group elements are located at the center only. surrounded by intermediate-mass elements which are mixed with oxygen.," As shown in Figure \ref{fig:comp}, iron group elements are located at the center only, surrounded by intermediate-mass elements which are mixed with oxygen."772 Most carbon is found in the outermost parts. but there is some mixing with intermediate-mass elements.," Most carbon is found in the outermost parts, but there is some mixing with intermediate-mass elements."773 It is important to note that both. carbon and oxygen. are at different locations than the iron. group elements. i.e. there is no mixing of unburned material with iron group elements.," It is important to note that both, carbon and oxygen, are at different locations than the iron group elements, i.e. there is no mixing of unburned material with iron group elements."774 There is also a clear difference between the distribution in the x—v--plane. which is the plane of rotation. and the distribution along the z--axis.," There is also a clear difference between the distribution in the -plane, which is the plane of rotation, and the distribution along the -axis."775 Along the c--axis. the iron group elements are spread out much more. which means that in this direction there are iron group elements at higher velocities than within the 1-v--plane.," Along the -axis, the iron group elements are spread out much more, which means that in this direction there are iron group elements at higher velocities than within the -plane."776 The opposite trend is observed for oxygen and intermediate mass elements. which reachfurther out in the plane of rotation than perpendicular to it.," The opposite trend is observed for oxygen and intermediate mass elements, which reachfurther out in the plane of rotation than perpendicular to it."777 This is a result of the density structure of the merged object through which the detonation propagates., This is a result of the density structure of the merged object through which the detonation propagates.778 Consequently. lightcurves and spectra of this explosion are expected to show considerable viewing angle dependence ?)..," Consequently, lightcurves and spectra of this explosion are expected to show considerable viewing angle dependence \citep[for viewing779 angle dependent lightcurves see][]{pakmor2010a}. ."780"Thus, where the second term of equation (42)) dominates the convective flux.","Thus, where the second term of equation \ref{eqn:qconv}) ) dominates the convective flux."781 Figure 6 shows these two heat fluxes normalized to the initial fiducial heat flux., Figure \ref{fig:A2-heatflux} shows these two heat fluxes normalized to the initial fiducial heat flux.782" The mass advection heat flux is enhanced due to the inability of the dark matter profile to adjust to changes in the ICM structure, so it is not a physically important quantity."," The mass advection heat flux is enhanced due to the inability of the dark matter profile to adjust to changes in the ICM structure, so it is not a physically important quantity."783" What is interesting, however, is that the convective heat flux is quite small."," What is interesting, however, is that the convective heat flux is quite small."784" In fact, it is roughly four orders of magnitude smaller than the conductive heat flux."," In fact, it is roughly four orders of magnitude smaller than the conductive heat flux."785 The small convective heat flux seems to result from two principal causes., The small convective heat flux seems to result from two principal causes.786" First, unlike in solar convection, there is a second channel open for energy flow, namely directly through conduction which is quite efficient."," First, unlike in solar convection, there is a second channel open for energy flow, namely directly through conduction which is quite efficient."787" Second, as the magnetic fields become more radial, the buoyant driving is significantly reduced, reducing the baseline convective motions."," Second, as the magnetic fields become more radial, the buoyant driving is significantly reduced, reducing the baseline convective motions."788 At this point it is also appropriate to discuss run A8., At this point it is also appropriate to discuss run A8.789" Run Αδ is chosen to be initialized at the MTI stability boundary for strong magnetic fields, namely kvA~Omax, for wavenumbers corresponding to scales slightly smaller than the cluster radius."," Run A8 is chosen to be initialized at the MTI stability boundary for strong magnetic fields, namely $kv_A \sim \sigma_{\textrm{max}}$, for wavenumbers corresponding to scales slightly smaller than the cluster radius."790" In this limit, all but the largest scale modes are suppressed entirely by magnetic tension."," In this limit, all but the largest scale modes are suppressed entirely by magnetic tension."791 There is both theoretical motivation and numerical evidence that the conductive flux increases as the initial magnetic field increases., There is both theoretical motivation and numerical evidence that the conductive flux increases as the initial magnetic field increases.792" We do indeed find that for run A8, the convective flux is within an order of magnitude of the conductive flux; however, this phase lasts only a very short time (less than 500 Myr) and the net convective flux never is larger than 0.035Q."," We do indeed find that for run A8, the convective flux is within an order of magnitude of the conductive flux; however, this phase lasts only a very short time (less than 500 Myr) and the net convective flux never is larger than $0.035 \,\widetilde{Q}$."793" In addition, for this run, the total heat flux carried is considerably smaller than the other runs."," In addition, for this run, the total heat flux carried is considerably smaller than the other runs."794" Thus, we conclude that conduction is the dominant heat transport mechanism and that convective heat transport can never be larger than a small fraction of the Spitzer conductivity for MTI-driven turbulence in galaxy clusters."," Thus, we conclude that conduction is the dominant heat transport mechanism and that convective heat transport can never be larger than a small fraction of the Spitzer conductivity for MTI-driven turbulence in galaxy clusters."795" In order to emphasize the nature of the MTI as a truly convective instability, we have performed simulations with purely isotropic conductivity."," In order to emphasize the nature of the MTI as a truly convective instability, we have performed simulations with purely isotropic conductivity."796 A long-favored approach to understanding thermal conduction in galaxy clusters has been to assume a tangled magnetic field geometry passively determined by hydrodynamic turbulence., A long-favored approach to understanding thermal conduction in galaxy clusters has been to assume a tangled magnetic field geometry passively determined by hydrodynamic turbulence.797 This tangled geometry then provides an effective thermal conductivity that is a specified fraction of Spitzer (???)..," This tangled geometry then provides an effective thermal conductivity that is a specified fraction of Spitzer \citep{rr78, cc98, nm01}."798" A major point of this paper is that such an approach is incorrect since thermal conduction along the magnetic field self-consistently modifies the magnetic field geometry, resulting in an that evolves with time and is not consistent with these fgpitzertheories."," A major point of this paper is that such an approach is incorrect since thermal conduction along the magnetic field self-consistently modifies the magnetic field geometry, resulting in an $f_{\textrm{Spitzer}}$ that evolves with time and is not consistent with these theories."799" Thus, here we present simulations I1 and I2 which have a thermal diffusivity fixed at or of the Spitzer value, respectively, for comparison purposes."," Thus, here we present simulations I1 and I2 which have a thermal diffusivity fixed at or of the Spitzer value, respectively, for comparison purposes."800" The temperature profiles of these two runs evolve in ways quite similar to that of run A2 and A3, reaching fairly isothermal in 5-7 Gyr."," The temperature profiles of these two runs evolve in ways quite similar to that of run A2 and A3, reaching fairly isothermal in 5-7 Gyr."801 A much more illustrative difference is found by considering the time evolution of the two runs., A much more illustrative difference is found by considering the time evolution of the two runs.802" The runs are initialized with a Mach 0.01 velocity perturbation that decays in time quite quickly, peaking at 1 Gyr or less, as opposed to 7 Gyr for run A2."," The runs are initialized with a Mach 0.01 velocity perturbation that decays in time quite quickly, peaking at 1 Gyr or less, as opposed to 7 Gyr for run A2."803" Without the MTI, the magnetic field is barely amplified at all, 6(B2)zx1.3 for both cases."," Without the MTI, the magnetic field is barely amplified at all, $\delta\langle B^2\rangle \approx 1.3$ for both cases."804" Finally, the magnetic field geometry is somewhat rearranged by the cluster motions, but does not even reach a geometrically isotropic state, with reaching à maximum of 0.19."," Finally, the magnetic field geometry is somewhat rearranged by the cluster motions, but does not even reach a geometrically isotropic state, with reaching a maximum of 0.19."805" Of course, we started in à state of pure azimuthal fields, so the key message here is that a field that started isotropically distributed would remain isotropically distributed—there is no force here pushing the field to be radially biased as with the MTI."," Of course, we started in a state of pure azimuthal fields, so the key message here is that a field that started isotropically distributed would remain isotropically distributed—there is no force here pushing the field to be radially biased as with the MTI."806" Clearly the comparison between these isotropic conduction calculations and the preceeding anisotropic heat conduction calculation shows that the MTI plays a significant role in setting a self-consistent solution for magnetic field amplification, convective velocities, and for the magnetic field geometry of a cluster."," Clearly, the comparison between these isotropic conduction calculations and the preceeding anisotropic heat conduction calculation shows that the MTI plays a significant role in setting a self-consistent solution for magnetic field amplification, convective velocities, and for the magnetic field geometry of a cluster."807" We now turn to the case of tangled magnetic fields, examined in run A3."," We now turn to the case of tangled magnetic fields, examined in run A3."808" This run has the exact same set-up as the previously discussed runs, except the intial magnetic field is a turbulent magnetic field distributed according to the Kolmogorov scaling law."," This run has the exact same set-up as the previously discussed runs, except the intial magnetic field is a turbulent magnetic field distributed according to the Kolmogorov scaling law."809" Thus, in this initial case, there is a zeroth order anisotropic heat flux that exists independent of the MTT's evolution."," Thus, in this initial case, there is a zeroth order anisotropic heat flux that exists independent of the MTI's evolution."810" Figure 7 shows the fairly rapid evolution of the temperature profile for the tangled field, reflecting this zeroth order"," Figure \ref{fig:A3-temp}811 shows the fairly rapid evolution of the temperature profile for the tangled field, reflecting this zeroth order"812As far as the exercise in this paper is concerned. it is more important to emphasise. again. the discrepancy between lines in Fie. 6.,"As far as the exercise in this paper is concerned, it is more important to emphasise, again, the discrepancy between lines in Fig. \ref{tz},"813 which calls into question how well the evolution in specific star formation rates hasreally boon captured by surveys thus far., which calls into question how well the evolution in specific star formation rates has been captured by surveys thus far.814 Very cilferent trends of stellar mass to star formation rate have been attributed. to the same observational sample2010)., Very different trends of stellar mass to star formation rate have been attributed to the same observational sample.815. In this paper. we have carried out a more rigorous investigation into the origin of these two key estimated. physical quantities.," In this paper, we have carried out a more rigorous investigation into the origin of these two key estimated physical quantities."816 A simple argument. based. on. hierarchical galaxy formation theory. was presented to understand why a strong trend might exist between these two quantities. anc that such a relationship can result from the self-similar nature of galaxy assembly. independently: of star. formation or [cedback processes.," A simple argument, based on hierarchical galaxy formation theory, was presented to understand why a strong trend might exist between these two quantities, and that such a relationship can result from the self-similar nature of galaxy assembly, independently of star formation or feedback processes."817 Using a mocel galaxy population as a guide. we have shown that. due to the combined. clleets of selection bias and physical scatter in the relations between observable and physical properties. this underlving trend can be easily nusrepresented.," Using a model galaxy population as a guide, we have shown that, due to the combined effects of selection bias and physical scatter in the relations between observable and physical properties, this underlying trend can be easily misrepresented."818" ""There can also be a large discrepancy between inferred and true parameter values.", There can also be a large discrepancy between inferred and true parameter values.819 These results highlight the importance of using realistic. physical galaxy formation models to guide the interpretation of high-redshift surveys.," These results highlight the importance of using realistic, physical galaxy formation models to guide the interpretation of high-redshift surveys."820 By subjecting mocdel-sgenerated galaxy populations to the same analysis as the real data. observation can be compared with competing theories on an even footing.," By subjecting model-generated galaxy populations to the same analysis as the real data, observation can be compared with competing theories on an even footing."821 In this wav. new surveys can Lead to more incisive quantitative conclusions about the true underlving galaxy population.," In this way, new surveys can lead to more incisive quantitative conclusions about the true underlying galaxy population."822 The authors would like to. thank Andrew Benson and Claudia Lagos for their helpful comments., The authors would like to thank Andrew Benson and Claudia Lagos for their helpful comments.823 CSE acknowledges a Roval Society Wolfson Research Merit Award and SAIC acknowledges the support of the Leverhulme Trust. Research Fellowship., CSF acknowledges a Royal Society Wolfson Research Merit Award and SMC acknowledges the support of the Leverhulme Trust Research Fellowship.824 DPS acknowledges support from an STEC postdoctoral research fellowship., DPS acknowledges support from an STFC postdoctoral research fellowship.825 This work was supported by an STEC rolling grant to the Institute for Computational Cosmology., This work was supported by an STFC rolling grant to the Institute for Computational Cosmology.826 The mock galaxy population in the figures in this article was generated using a version of the semi-analvtic model which is currenth under development., The mock galaxy population in the figures in this article was generated using a version of the semi-analytic model which is currently under development.827 The goal of this new version is to combine the most realistic aspects of the two previously published versions2006).. whilst achieving a better match. than either of these. to current observational constraints.," The goal of this new version is to combine the most realistic aspects of the two previously published versions, whilst achieving a better match, than either of these, to current observational constraints."828 The development model is most closely related to the model published by(2006).. but uses parameters [or star. formation and feedback.et that are more realistic: closer to those favoured by(2005).," The development model is most closely related to the model published by, but uses parameters for star formation and feedback that are more realistic; closer to those favoured by."829. A list of parameter changes appears in table Al., A list of parameter changes appears in table \ref{parameters}.830" The parameters 7, and a, apply to star formation rate c. as follows: These. and all other parameters. are as defined in(2006)... and references therein."," The parameters $\tau_\star$ and $\alpha_\star$ apply to star formation rate $\psi$, as follows: = These, and all other parameters, are as defined in, and references therein."831 Other changes are the distribution of. orbital parameters. which has been updated to. follow(2005).. and the treatment of the cooling of hot halo gas. which now follows(2010).," Other changes are the distribution of orbital parameters, which has been updated to follow, and the treatment of the cooling of hot halo gas, which now follows."832. Full details of this version will appear in Lacey et al. (, Full details of this version will appear in Lacey et al. (8332011. in prep.),"2011, in prep.)."834 To explore this relationship. we return to the moclel galaxy population that was shown in Fig. 2..," To explore this relationship, we return to the model galaxy population that was shown in Fig. \ref{hist}."835 Consider. first. their star formation rate vs. rest-[rame UV magnitude. which is shown in the left main panel of Fig. Bl.," Consider, first, their star formation rate vs. rest-frame UV magnitude, which is shown in the left main panel of Fig. \ref{comp}."836 The relation assumed. in observational analysis is shown in the, The relation assumed in observational analysis is shown in the837The acceleration during the planar phase. ie.. before a shell doubles its radius. follows (Johnson&Melee1971:PanSari200G)::where / is the lab frame time (not to be confused with observer time) measured since breakout.,"The acceleration during the planar phase, i.e., before a shell doubles its radius, follows \citep{Johnson71,Pan06}:where $t$ is the lab frame time (not to be confused with observer time) measured since breakout."838 If the shell is optically thick at the end of acceleration aud acceleration ends during the planar phase then (he final Lorentz [actor is This relation is general as it does not depend on the exact densitv. profile. as long as there is a large energy reservoir behind the accelerating shell.," If the shell is optically thick at the end of acceleration and acceleration ends during the planar phase then the final Lorentz factor is This relation is general as it does not depend on the exact density profile, as long as there is a large energy reservoir behind the accelerating shell."839 The acceleration ends al Thus. more massive shells end their acceleration at earlier times and lower Lorentz factors.," The acceleration ends at Thus, more massive shells end their acceleration at earlier times and lower Lorentz factors."840" In (his paper we restrict our treatment (o cases where the shell ends its acceleration dung the planar phase. ie. /pj«/, where is the lab Iranmie time of transition between the planar and spherical phases."," In this paper we restrict our treatment to cases where the shell ends its acceleration during the planar phase, i.e., $t_{f,0}<t_s$ where is the lab frame time of transition between the planar and spherical phases."841" Since optical depth of the shell is of order unity (oe. &ppgzo£8 1): where AZ, and A, are (he progenitor star mass and radius in units of à solar masses ancl radii respectively."," Since pre-shocked optical depth of the shell is of order unity (i.e., $\kappa_T \rho_0 z_0 \approx 1$ ): where $M_x$ and $R_x$ are the progenitor star mass and radius in units of $x$ solar masses and radii respectively."842 Thus. there is a critical Lorentz factor. 590. for which /po= {ει The corresponding critical final Lorentz factor of the shell is:," Thus, there is a critical Lorentz factor, $\g_{0,s}$, for which $t_{f,0}=t_s$ The corresponding critical final Lorentz factor of the shell is:"843projection of the fundamental plane.,projection of the fundamental plane.844 Galaxies following the virial relation would fall on a line with a slope of one., Galaxies following the virial relation would fall on a line with a slope of one.845 Observed galaxies do not fall on the expected virial relation., Observed galaxies do not fall on the expected virial relation.846 This variance from virial scaling is the so-called tilt of the fundamental plane., This variance from virial scaling is the so-called tilt of the fundamental plane.847 In the third panel of Figure 2(a) we plot Mstar versus σ΄R., In the third panel of Figure \ref{fig:toyGasConst} we plot $M_{\rm star}$ versus $\sigma^2R$.848 The line plotted has a slope of one and therefore follows a virial scaling., The line plotted has a slope of one and therefore follows a virial scaling.849 The remnants from each gas fraction set fall on the same virial scaling., The remnants from each gas fraction set fall on the same virial scaling.850 This demonstrates that our model produces no tilt in the fundamental plane for remnants when the progenitors’ properties are scaled with the mass., This demonstrates that our model produces no tilt in the fundamental plane for remnants when the progenitors' properties are scaled with the mass.851" However, while each gas fraction series follows the virial scaling, there is a shift between cases with different"," However, while each gas fraction series follows the virial scaling, there is a shift between cases with different"852We have investigated the radio allerglows of GRBs and IINe at high redshifts. analvzing their detectability as well as their possible usefulness for 21 em absorption measurenients.,"We have investigated the radio afterglows of GRBs and HNe at high redshifts, analyzing their detectability as well as their possible usefulness for 21 cm absorption measurements."853 We explored several possible sets of model parameters which are plausible at high redshifts (see Table 1))., We explored several possible sets of model parameters which are plausible at high redshifts (see Table \ref{tab:model}) ).854 Our source models include the effects associated with reverse as well as forward shock emission. jel breaks and sidewavs expansion of the jet. and (he transition [rom a relativistic to non-relativistic expansion regime.," Our source models include the effects associated with reverse as well as forward shock emission, jet breaks and sideways expansion of the jet, and the transition from a relativistic to non-relativistic expansion regime."855 Our results indicate that standard GRD radio afterglows may be detected: with the current. VLA up to ze30. and energetic LIN afterelows maa be detected up to z~20. a high radio frequencies 5 Gllz.," Our results indicate that standard GRB radio afterglows may be detected with the current VLA up to $z \sim 30$, and energetic HN afterglows may be detected up to $z \sim 20$, at high radio frequencies $\sim 5$ GHz."856 The proposed SIVA should be easily able to detect the GRB and ΗΝ alterglows bevoud 2~30 at frequencies 5 GIIz., The proposed SKA should be easily able to detect the GRB and HN afterglows beyond $z \sim 30$ at frequencies $\sim 5$ GHz.857 Other possible effects which max be expected at high redshift. such as a larger isotropic equivalent energv. higher external density. longer duration GRB and the possibility of magnetized ejecta. can all contribute to enhance the detectability of afterglows at ~5 GIIz.," Other possible effects which may be expected at high redshift, such as a larger isotropic equivalent energy, higher external density, longer duration GRB and the possibility of magnetized ejecta, can all contribute to enhance the detectability of afterglows at $\sim 5$ GHz."858 We find that it is difficult to detect the 21 cm absorption line due to a dilfuse neutral ICM. even with the SIXA.," We find that it is difficult to detect the 21 cm absorption line due to a diffuse neutral IGM, even with the SKA."859 A possible exception may be in the case of a very high isotropic equivalent ΗΝ energy E~10?! ere. if it occurs in a low ambient densitv n£1 *a( an epoch with low spin temperature Ty<0.57t445.," A possible exception may be in the case of a very high isotropic equivalent HN energy $E\sim 10^{54}$ erg, if it occurs in a low ambient density $n \siml 1$ $^{-3}$ at an epoch with low spin temperature $T_{S} \siml 0.5 T_{\rm CMB}$."860 This case is very hypothetical. but could be plausible under some conditions.," This case is very hypothetical, but could be plausible under some conditions."861 The 21 em absorption [rom collapsed gas clouds along the line of sight. e.g.. [from minihalos and galactic disks at 26 and damped Lya systems al 2S6. is in principle detectable in the energetic IIN case. and marginally detectable in the GRB alterglow cases considered. if the 21 em optical depth of the clouds is larger than unity and the ambient gas density of the GRBs is low. nZLem 7.," The 21 cm absorption from collapsed gas clouds along the line of sight, e.g., from minihalos and galactic disks at $z \simg 6$ and damped $\alpha$ systems at $z \siml 6$, is in principle detectable in the energetic HN case, and marginally detectable in the GRB afterglow cases considered, if the 21 cm optical depth of the clouds is larger than unity and the ambient gas density of the GRBs is low, $n \siml 1$ $^{-3}$."862 However. such high optical depth gas clouds are rare. and (their detection is not practical given the long integration (times needed.," However, such high optical depth gas clouds are rare, and their detection is not practical given the long integration times needed."863 This indicates that if 21 cm absorption is detected in a narrow range of frequencies. it is likely {ο originate in the host galaxy of the GRB or UN with a high 21 em optical depth. e.g.. in an edge-on disk. star-forming galaxy or protogalaxy. and may thus provide a direct measurement of the redshift of the host.," This indicates that if 21 cm absorption is detected in a narrow range of frequencies, it is likely to originate in the host galaxy of the GRB or HN with a high 21 cm optical depth, e.g., in an edge-on disk, star-forming galaxy or protogalaxy, and may thus provide a direct measurement of the redshift of the host."864 We (hank ALJ. Rees and S. Ixobavashi for useful discussions., We thank M.J. Rees and S. Kobayashi for useful discussions.865 This work was supported in part bv the Eberly Research Funds of Penn State aud by (he Center for Gravitational Wave Physics under grants PIIY-01-14375 (INI). NASA NAGS5-13286. NSF AST 0098416 and the Monell Foundation.," This work was supported in part by the Eberly Research Funds of Penn State and by the Center for Gravitational Wave Physics under grants PHY-01-14375 (KI), NASA NAG5-13286, NSF AST 0098416 and the Monell Foundation."866"search oscillation modes with angular degree (=0, 1. 2 and in the domain of frequencies defined by an interval around ως (Eq.(10)ofKjeldsen&Bedding1995).","search oscillation modes with angular degree $\ell=0$, 1, 2 and in the domain of frequencies defined by an interval around $\nu_{\rm max}$ \citep[Eq. (10) of][]{KB95}."867. The width of the solar-like frequency domain is taken to be larger than the difference between the acoustic cutoff frequency in the stellar atmosphere and 1..., The width of the solar-like frequency domain is taken to be larger than the difference between the acoustic cutoff frequency in the stellar atmosphere and $\nu_{\rm max}$.868" The properties of oscillation modes depend on the behaviour of the (CV) and Lamb (5,) frequencies.", The properties of oscillation modes depend on the behaviour of the $N$ ) and Lamb $S_{\ell}$ ) frequencies.869" In red-giant models, Vv reaches huge values in the central regions and therefore the frequency of gravity modes (g-modes) and their number by frequency interval (1,) Increase with respect to main sequence models."," In red-giant models, $N$ reaches huge values in the central regions and therefore the frequency of gravity modes (g-modes) and their number by frequency interval $n_g$ ) increase with respect to main sequence models."870 On the other hand. the low mean density makes the frequency of pressure modes (p-modes) to decrease.," On the other hand, the low mean density makes the frequency of pressure modes (p-modes) to decrease."871" All that leads to an oscillation spectrum for red-giants where, in addition to radial modes, one finds a large number of non-radial modes with mixed e-p properties."," All that leads to an oscillation spectrum for red-giants where, in addition to radial modes, one finds a large number of non-radial modes with mixed g-p properties."872" The dominant character of these non-radial modes depends on the separation between gravity and acoustic cavities, and may be estimated from the value of the normalized mode inertia (77) (seee.g.Christensen-Dalsgaard2004,andreferencestherein).."," The dominant character of these non-radial modes depends on the separation between gravity and acoustic cavities, and may be estimated from the value of the normalized mode inertia $E$ ) \citep[see e.g.][ and references therein]{JCD04}."873" Therefore. some non-radial modes may be well trapped in the acoustic cavity and behave as presenting a mode inertia close to that of radial modes, while modes with strong mixed σ-ρ character have larger 77 value."," Therefore, some non-radial modes may be well trapped in the acoustic cavity and behave as p-modes presenting a mode inertia close to that of radial modes, while modes with strong mixed g-p character have larger $E$ value."874" Hereatter, we will use the term p-modes in quotation marks to refer to mixed modes with a dominant p-character."," Hereafter, we will use the term p-modes in quotation marks to refer to mixed modes with a dominant p-character."875 In Fig., In Fig.876" | we present, in top panels, the /=1, 2 propagation diagrams for a RGB 1.5 mmodel (left) and for a core He-burning (He-B) model of ((right)."," \ref{label_prop} we present, in top panels, the $\ell=1$, 2 propagation diagrams for a RGB 1.5 model (left) and for a core He-burning (He-B) model of (right)."877 In the bottom panels we plot the variation of the mode inertia with frequency for radial and non-radial modes (f=1. 2).," In the bottom panels we plot the variation of the mode inertia with frequency for radial and non-radial modes $\ell=1,\,2$ )."878" As mentioned above, the RGB model is ten times more centrally condensed than the He-B one."," As mentioned above, the RGB model is ten times more centrally condensed than the He-B one."879 The huge difference in density between the central region and the convective envelope entails a high potential barrier between the acoustic and the gravity cavities reducing the interaction between p and σ modes., The huge difference in density between the central region and the convective envelope entails a high potential barrier between the acoustic and the gravity cavities reducing the interaction between p and g modes.880" As a consequence, we find for RGB models that (=| modes with νι£zFy) are quite regularly spaced in frequency."," As a consequence, we find for RGB models that $\ell=1$ modes with $E_{\ell=1}\approx E_{\ell=0}$ are quite regularly spaced in frequency."881" For He-B ones, the coupling between these cavities is more important and /=1 modes are mixed modes with Ey."," For He-B ones, the coupling between these cavities is more important and $\ell=1$ modes are mixed modes with $E_{\ell=1}> E_{\ell=0}$ ."882" Nevertheless, 77,4 presents still à minimum value for modes between two consecutive radial ones showing a somewhat regular pattern."," Nevertheless, $E_{\ell=1}$ presents still a minimum value for modes between two consecutive radial ones showing a somewhat regular pattern."883" Even if the F value is larger than that corresponding to radial modes we can still consider those modes, based on the value of E, as observable ""p-modes""."," Even if the $E$ value is larger than that corresponding to radial modes we can still consider those modes, based on the value of $E$, as observable “p-modes”."884" For (=2 modes, the coupling between the g- and p-cavities is smaller than for (=1 and hence the trapping more efficient."," For $\ell=2$ modes, the coupling between the g- and p-cavities is smaller than for $\ell=1$ and hence the trapping more efficient."885" Therefore, independently of the central condensation of the model, à spectrum of regularly spaced /=2 ""p-modes"" with yo.zzFy) 1s expected."," Therefore, independently of the central condensation of the model, a spectrum of regularly spaced $\ell=2$ “p-modes” with $E_{\ell=2}\approx E_{\ell=0}$ is expected."886" Finally, note that the turning points for acoustic modes (/p, defined as the point where μις= 5,) are inside the convective envelope for the RGB model and in the radiative region for the He-B one."," Finally, note that the turning points for acoustic modes $tp_\ell$ defined as the point where $\nu_{\rm max}=S_{\ell}$ ) are inside the convective envelope for the RGB model and in the radiative region for the He-B one."887Quantities related to the How are unchanged by the averaging procedure since the superlluid How velocity is independent of whether vortices are pinned or not.,Quantities related to the flow are unchanged by the averaging procedure since the superfluid flow velocity is independent of whether vortices are pinned or not.888 The factors of f; in eq. (49)), The factors of $f_v$ in eq. \ref{fave}) )889 account for the fact that only the motion of the translating vortex segments contributes to the mutual frietion (see. also. 2)).," account for the fact that only the motion of the translating vortex segments contributes to the mutual friction (see, also, \citealt{jahanmiri06}) )."890 The value of f is unimportant for the following estimates., The value of $f_v$ is unimportant for the following estimates.891 The force of eq. (47)).," The force of eq. \ref{fagain}) ),"892 which is appropriate for vortex creep. must equal the average force (f/f). giving the following relationships: We now use estimates of oyαυ to obtain the ratio 3/a.," which is appropriate for vortex creep, must equal the average force $\langle{\fbf_0}/\rho\rangle$, giving the following relationships: We now use estimates of $\beta_0/\alpha_0$ to obtain the ratio $\beta/\alpha$."893 The dominant drag process on unpinned vortex segments considered so far arises from the excitation of Kelvin modes as the vortex moves past nuclei., The dominant drag process on unpinned vortex segments considered so far arises from the excitation of Kelvin modes as the vortex moves past nuclei.894 Calculations of dissipation bv Ixelvin phonon production on a long vortex with periodic »oundary conditions for ey~LO geive typical values of ο=0.1 and ag~1 (2).," Calculations of dissipation by Kelvin phonon production on a long vortex with periodic boundary conditions for $v_0\sim 10^7$ give typical values of $\beta_0/\alpha_0= 0.1$ and $\alpha_0\sim8951$ \citep{eb92}."896 Pinning occurs [or eg107.Land «οαυ is likely to be significantly smallerin this velocity regime due to strong suppression of Ixelvin phonon production (?)..," Pinning occurs for $v_0\lap 10^5$, and $\beta_0/\alpha_0$ is likely to be significantly smaller in this velocity regime due to strong suppression of Kelvin phonon production \citep{jones92}."897 Vortex ereep is therefore a low-clrag »ocess if Ixelvin phonon production is the dominant dissipative mechanism., Vortex creep is therefore a low-drag process if Kelvin phonon production is the dominant dissipative mechanism.898 We fix τα=0.1 for illustration in the following. which we consider to be an upper limit: we expect tvpical values to be smaller.," We fix $\beta/\alpha=0.1$ for illustration in the following, which we consider to be an upper limit; we expect typical values to be smaller."899 We nowestimate 3., We nowestimate $\beta$.900 We adopt polar coordinates (7.0.2). with the unperturbecl vorticity along ὁ and the unperturbed low vy along ὦ. and take the unperturbed flow and vortex velocity field. to be axisvmmoetric.," We adopt polar coordinates $(r,\phi,z)$, with the unperturbed vorticity along $\hat{z}$ and the unperturbed flow $\vbf_0$ along $\hat{\phi}$, and take the unperturbed flow and vortex velocity field to be axisymmetric."901 In the rotating frame. the unperturbed vortex velocity from eq. (19))," In the rotating frame, the unperturbed vortex velocity from eq. \ref{vv0}) )"902 is where 5 is the average direction of vortex motion., is where $\hat{n}$ is the average direction of vortex motion.903 For steady spin down of the star. the inner crust superfIuid and the crust are spinning down atthe same rate for a local differential velocity ey.," For steady spin down of the star, the inner crust superfluid and the crust are spinning down atthe same rate for a local differential velocity $v_0$."904 Phe creep velocity in this steady state is related to the spin-down rate by (72) where © is the spin rate of the supertluicl. Oy is the observed spin downrate of the crust. and rr is approximately the stellar racius 7.," The creep velocity in this steady state is related to the spin-down rate by \citep{alpar_etal84,leb93} where $\Omega$ is the spin rate of the superfluid, $\dot{\Omega}_0$ is the observed spin downrate of the crust, and $r$ is approximately the stellar radius $R$."905 We arrive at the estimate where QcQy is assumed. and fase=Qu/2|Qu| is spin-down age.," We arrive at the estimate where $\Omega\simeq\Omega_0$ is assumed, and $t_{\rm age}\equiv906\Omega_0/2\vert\dot{\Omega}_0\vert$ is the spin-down age."907 Eq. (53)).," Eq. \ref{ss}) ),"908 with 3= θα. gives the fiducial value 121074.," with $\beta=0.1\alpha$ , gives the fiducial value $\alpha\beta=10^{-21}$."909 For this value. we deduce fo~(odfoagdu)?7be10)1 that is. most of the vortex length is pinned at any instant.," For this value, we deduce $f_v\sim (\alpha\beta/\alpha_0\beta_0)^{1/2}\sim 10^{-11}$, that is, most of the vortex length is pinned at any instant."910 The unperturbed vortex creep speed. from eq. (513).," The unperturbed vortex creep speed, from eq. \ref{vv0again}) ),"911 is ~ary107 ivo. justifving the neglect of Or.Οἱ compared to vy in the stability analysis.," is $\sim\alpha\, v_0\sim 10^{-5}$ $<<v_0$ , justifying the neglect of $\partial\rbf_{v0}/\partial t$ compared to $\vbf_0$ in the stability analysis."912 We can now proeced with estimates of the instability length scale and. growth rate., We can now proceed with estimates of the instability length scale and growth rate.913 For 3<a and 6—tan.1(//2) in eq. (42)).," For $\beta<\alpha$ and $\theta=\tan^{-1}(\sqrt{2})$ in eq. \ref{kc}) ),"914 the critical wavenumber is corresponding to a wavelength A=2x/kz10 m. For &cc Ay. the erowth rate from eq. (43))," the critical wavenumber is corresponding to a wavelength $\lambda=2\pi/k\simeq 10$ m. For $k>>k_c$ , the growth rate from eq. \ref{highk}) )"915 is The hydrodynamic treatment is restricted to keg<<ο., is The hydrodynamic treatment is restricted to $kc_T<<\Omega$.916 To est high the growth mue could be. we consider a maximum wavenumber defined by ομως=0.10. where erc10imate1(0/100.pwrads1957.1.," To estimate how high the growth rate could be, we consider a maximum wavenumber defined by $c_Tk_{\rm max}=0.1\,\Omega$, where $c_T\simeq 10^{-1}\, (\Omega/100\mbox{ \rads})^{1/2}$."917 The growth rate at this waventunber. from eq. (43)).," The growth rate at this wavenumber, from eq. \ref{highk}) ),"918 is For ο=100 rad J|. the corresponding wavenumber is Aac100 1," is For $\Omega=100$ rad $^{-1}$, the corresponding wavenumber is $k_{\rm max}\simeq 100$ $^{-1}$."919 Ίσα. (56)), Eq. \ref{highsigma}) )920 does not represent a physical limit. but only the restrictions of the hydrodynamic treatment: the instability could continue to exist also for wavenumbers in the regime ker> Q.," does not represent a physical limit, but only the restrictions of the hydrodynamic treatment; the instability could continue to exist also for wavenumbers in the regime $kc_T>\Omega$ ."921 If vortex creep is in the strongly-clamiped regime 7 a. contrary to the estimateshere. there is still abroad. window for instability.," If vortex creep is in the strongly-damped regime $\beta>>\alpha$ , contrary to the estimateshere, there is still abroad window for instability."922 RequiringAe«Aue gives and the star will be unstable at some wavenumber that is consistent with the hvdrodyvnamic regime hep<< ο., Requiring$k_c<k_{\rm max}$ gives and the star will be unstable at some wavenumber that is consistent with the hydrodynamic regime $kc_T<<\Omega$ .923the sky well removed from the Galactic equator. where the celestial diffuse euission is particularly intense.,"the sky well removed from the Galactic equator, where the celestial diffuse emission is particularly intense."924 The denoising results reported iu Section 5Γ use a data cube obtained according to this simulation scenario., The denoising results reported in Section \ref{sect_glast_exp} use a data cube obtained according to this simulation scenario.925 Tn this section. we review the inethod(2). restricted to the Isotropic Uuclecimated Wavelet Transform (IIUNT).," In this section, we review the method, restricted to the Isotropic Undecimated Wavelet Transform (IUWT)."926 Indeed. the call Use other transforms such the standard uudecimated wavelet transform. the ridgelet or the curvelet transforms: see(7).," Indeed, the can use other transforms such the standard undecimated wavelet transform, the ridgelet or the curvelet transforms; see."927. In our specific casehere.. oulv the is of interest.," In our specific case, only the is of interest."928 Given X a sequence of n» indepeudeut Poisson random variables Nj.f=1.5.20. each of mean A;. y=arVh]X;j the filkered process obtained by convolving the sequence X with a discrete filter A.," Given $\fX$ a sequence of $n$ independent Poisson random variables $X_i, i=1,\cdots,n,$ each of mean $\lambda_i$, $Y_i = \sum_{j=1}^n h[j]X_{i-j}$ the filtered process obtained by convolving the sequence $\fX$ with a discrete filter $h$ ."929" Y. denotes auv one of the τα, and 74=Σω. for kh—123.6."," $Y$ denotes any one of the $Y_i$ 's, and $\tau_k = \sum_i (h[i])^k$ for $k=1,2,\cdots$."930 If)=à. then the Auscombe of Y; (hence Vy) acts as if the data arose roni a Gaussian white noise with uuit variance. under he asstunption that the intensity A; is large.," If $h=\delta$, then the Anscombe of $Y_i$ (hence $X_i)$ acts as if the data arose from a Gaussian white noise with unit variance, under the assumption that the intensity $\lambda_i$ is large."931 This is why he Anscombe VST performs poorly in low-count scttines., This is why the Anscombe VST performs poorly in low-count settings.932 But. if the filterPh acts as an “averaging” kernel (auore eenerallv a low-pass filter). one can reasonably expect that stabilizing 3; would be more beneficial. since the signal ratio neasured at the output of / is expected to )o higher.," But, if the filter$h$ acts as an “averaging” kernel (more generally a low-pass filter), one can reasonably expect that stabilizing $Y_i$ would be more beneficial, since the signal-to-noise ratio measured at the output of $h$ is expected to be higher."933 Using a local homogeneity assuniptiou. Le. Ap;=A or all; within the support of /.it has been shown that or a non-negative filter 5. the transform Z=by|ο with b>0 and e>0 defined as is a second order accurate variance stabilization transform. with asviuptotic wit variance.," Using a local homogeneity assumption, i.e. $\lambda_{i-j}=\lambda$ for all $j$ within the support of $h$, it has been shown that for a non-negative filter $h$, the transform $Z = b \sqrt{Y + c}$ with $b > 0$ and $c >0$ defined as c = - b = is a second order accurate variance stabilization transform, with asymptotic unit variance."934" By second-order accurate, we mean that the error term in the variance of the stabilized variableZ decreases rapidly as O(A3], From(1).."," By second-order accurate, we mean that the error term in the variance of the stabilized variable$Z$ decreases rapidly as $O(\lambda^{-2})$."935 if is obvious that when h=à. we obtain the classical Anscombe VST paramcters 5=2 and ο=3/8.," From, it is obvious that when $h=\delta$ , we obtain the classical Anscombe VST parameters $b=2$ and $c=3/8$ ."936 The authors in lave also proved that Z is asvinptotically distributed as a Gaussian variate with mean byf7A and unit variance., The authors in have also proved that $Z$ is asymptotically distributed as a Gaussian variate with mean $b \sqrt{\tau_1\lambda}$ and unit variance.937 A non-positive 5 with a negative c could also be considered: see for more details., A non-positive $h$ with a negative $c$ could also be considered; see for more details.938 Fig.2 shows the estimates of obtained from 9:10 Poisson noise realizatious of X. plotted as a function of the intensity A for both Auscombe (clashed-dotted). Ilaar-Fisz and our VST with the 2D Bs-Spline filter as a low-pass filter / (solid).," \ref{fig_msvst} shows the estimates of obtained from $2 \cdot 10^5$ Poisson noise realizations of $\fX$, plotted as a function of the intensity $\lambda$ for both Anscombe (dashed-dotted), Haar-Fisz and our VST with the 2D $B_3$ -Spline filter as a low-pass filter $h$ (solid)."939 The axviiptotie bounds (dots) (i.e. 1 for the variance aud VAfor the expectation) are also shown., The asymptotic bounds (dots) (i.e. $1$ for the variance and $\sqrt{\lambda}$for the expectation) are also shown.940 It can be secu hat for increasing iutensity. E|Z| aud Var|Z] approach he theoretical bounds at different rates depenudiug ou the VST used.," It can be seen that for increasing intensity, $\bE[Z]$ and $\var{Z}$ approach the theoretical bounds at different rates depending on the VST used."941 Quantitatively. Poisson. variables trausformed using the Anscombe VST can )o reasonably considered to τὸ unbiased and stabilized for Az10. using Ibuu-Fisz for Az l. and using out (after low-pass filtering with he chosen 7) for Az0.1.," Quantitatively, Poisson variables transformed using the Anscombe VST can be reasonably considered to be unbiased and stabilized for $\lambda \gtrapprox 10$, using Haar-Fisz for $\lambda \gtrapprox 1$ , and using out (after low-pass filtering with the chosen $h$ ) for $\lambda \gtrapprox 0.1$."942 The uudeciaated: wavelet transform (UWT) uses au analysis filter bauk (5.9) to decompose a signal ay iuto a cocficient set Woo={dy.....γα]. where dj is the wavelet (detail) coefficients. at scale jj and d; is the approximation cocficicuts at the coarsest resolution J.," The undecimated wavelet transform (UWT) uses an analysis filter bank $(h,g)$ to decompose a signal $a_0$ into a coefficient set $W = \{d_1, \dots, d_J, a_J\}$, where $d_j$ is the wavelet (detail) coefficients at scale $j$ and $a_J$ is the approximation coefficients at the coarsest resolution $J$."943" The passage from one resolution to the next one is obtained using the ""à trous” algorithin where Εξhl]if !/2/c Z aud ϐ otherwise. Alf}= h|Ff. aud ""« denotes discrete circular convolution."," The passage from one resolution to the next one is obtained using the “à trous” algorithm where $h^{\uparrow j}[l] = h[l]$if $l / 2^j \in \bZ$ and $0$ otherwise, $\bar{h}[l] = h[-l]$ , and $\star$ ” denotes discrete circular convolution."944 The reconstruction ds eiven by ZH=$[Mwayal|(ghtxej3d ., The reconstruction is given by $a_{j}[l] = \frac{1}{2}\left[ (\tilde{h}^{\uparrow j} \star a_{j+1})[l] + (\tilde{g}^{\uparrow j} \star w_{j+1})[l] \right]$ .945 The filterbank (fry.h.g) needs to satisfv the so-called exact reconstruction condition (?7).," The filterbank $(h,g,\tilde{h},\tilde{g})$ needs to satisfy the so-called exact reconstruction condition ."946. The Isotropic UWT (IUWT) uses the &lter bank (hy=àh.h-byà) wherefi is typically a sviunietric low-pass filter such as the B3-Splinefilter.," The Isotropic UWT (IUWT) uses the filter bank $(h,g=\delta-h,\tilde{h}=\delta,\tilde{g}=\delta)$ where$h$ is typically a symmetric low-pass filter such as the $B_3$ -Splinefilter."947 The, The948on scales smaller than this separation around the quasars we examine.,on scales smaller than this separation around the quasars we examine.949" At slightly larger radii we find no evidence for photometric objects being systematically excluded for spectroscopic follow-up, so we expect our analysis to be complete beyond this radius."," At slightly larger radii we find no evidence for photometric objects being systematically excluded for spectroscopic follow-up, so we expect our analysis to be complete beyond this radius."950" In addition to this photometric incompleteness, our analysis may also be limited by the spectroscopic resolution of the data."," In addition to this photometric incompleteness, our analysis may also be limited by the spectroscopic resolution of the data."951" The galaxy redshifts in the DEEP2 catalogue have a typical lo precision of o,=1.3x107* which corresponds a physical scale of ~400 h7'kpc for matter in the Hubble flow at z=1.", The galaxy redshifts in the DEEP2 catalogue have a typical $\sigma$ precision of $\sigma_{z}$ $\times10^{-4}$ which corresponds a physical scale of $\sim$ 400 $^{-1}$ kpc for matter in the Hubble flow at z=1.952" As such, detections of galaxies and absorbers that are physically associated at the same redshift may still have measured separations on the order of this limiting resolution."," As such, detections of galaxies and absorbers that are physically associated at the same redshift may still have measured separations on the order of this limiting resolution."953" To investigate any possible incompleteness in the data, we re-examined the spectra of all quasars with a projected separation of «200 h! kpc from DEEP2 galaxies."," To investigate any possible incompleteness in the data, we re-examined the spectra of all quasars with a projected separation of $<$ 200 $^{-1}$ kpc from DEEP2 galaxies."954" Stacking each of these 63 spectra in the rest-frame of the nearby galaxy revealed a 2c cumulative detection ofII, which grew in significance as we reduced the maximum projected separation allowed."," Stacking each of these 63 spectra in the rest-frame of the nearby galaxy revealed a $\sigma$ cumulative detection of, which grew in significance as we reduced the maximum projected separation allowed."955" After examining the individual spectra contributing to the stack, we determined that no absorbers with >3o significance were undetected by the automated algorithm."," After examining the individual spectra contributing to the stack, we determined that no absorbers with $>$ $\sigma$ significance were undetected by the automated algorithm."956" We also found that the bulk of the contribution to the absorption signal in the stacked spectrum could be attributed to the single, strong absorber with a 37 h! kpc separation from a DEEP2 galaxy (ID 16 in Table 1), which has been thoroughly discussed in Section 5."," We also found that the bulk of the contribution to the absorption signal in the stacked spectrum could be attributed to the single, strong absorber with a 37 $^{-1}$ kpc separation from a DEEP2 galaxy (ID 16 in Table 1), which has been thoroughly discussed in Section 5."957" The bias we measure for absorbers is consistent with that of the DEEP2 galaxies, although the measurement error is substantial due to the small size of the sample."," The bias we measure for absorbers is consistent with that of the DEEP2 galaxies, although the measurement error is substantial due to the small size of the sample."958 This basic agreement suggests that strong absorbers at z~1 reside in similar environments to those of the galaxies in the DEEP2 survey., This basic agreement suggests that strong absorbers at $\sim$ 1 reside in similar environments to those of the galaxies in the DEEP2 survey.959" Furthermore, the average halo mass we estimate, 1.5642:0.11x10? is consistent with observations at lower redshift."," Furthermore, the average halo mass we estimate, $\pm$ $\times10^{12}\Msun$ , is consistent with observations at lower redshift."960" LundgrenMo,etal.(2009) measured a typical halo mass of 1.812x101?h-!Mo for absorbers with W??798 at z=0.6, in agreement with Bouchéetal.(2006) and Gauthieretal. (2009)."," \citet{L09} measured a typical halo mass of $\pm^{4.2}_{1.6}\times10^{12}h^{-1}\Msun$ for absorbers with $_{r}^{\lambda2796}$ at z=0.6, in agreement with \citet{B06} and \citet{Gauthier09}."961". The consistency with our measurement at higher redshift suggests that the halo masses of absorbers evolve very little from z=1, though the errors remain large."," The consistency with our measurement at higher redshift suggests that the halo masses of absorbers evolve very little from z=1, though the errors remain large."962" It is also worth noting that the luminosity function of z~0.65 galaxies selected by absorption in nearby quasar sightlines has been shown to peak at Mpg--20 (Steideletal.1994),, which is consistent with the mean luminosity of the DEEP2 galaxies."," It is also worth noting that the luminosity function of $\sim$ 0.65 galaxies selected by absorption in nearby quasar sightlines has been shown to peak at $M_{B}$ =-20 \citep{Steideletal94}, which is consistent with the mean luminosity of the DEEP2 galaxies."963" While we have not analysed statistical sample of MgII-selected galaxies, it seems likelya that the z~1 absorbers not only trace the same environments as the DEEP2 galaxy sample but also similar types of galaxies (~ L5)."," While we have not analysed a statistical sample of -selected galaxies, it seems likely that the $\sim$ 1 absorbers not only trace the same environments as the DEEP2 galaxy sample but also similar types of galaxies $\sim L_{B}^{*}$ )."964" As measurements of the bias of have never been reported at this redshift, we can provide the first constraints on theoretical models describing the evolution of absorbers and their respective haloes."," As measurements of the bias of have never been reported at this redshift, we can provide the first constraints on theoretical models describing the evolution of absorbers and their respective haloes."965" Tinker&Chen(2008) produced a model for the halo occupation distribution of cold gas at z=0.6, where the bias ofII has been precisely measured from cross-correlations with LRGs."," \citet{TC08} produced a model for the halo occupation distribution of cold gas at z=0.6, where the bias of has been precisely measured from cross-correlations with LRGs."966" With this model, Tinker&Chen(2008) predict the probability of finding an absorber with an equivalent width W;. in a halo of mass Μη."," With this model, \citet{TC08} predict the probability of finding an absorber with an equivalent width $_{r}$ in a halo of mass $_{h}$."967" In doing so, they demonstrate that the observed anti-correlation of and bias at z=0.6 may be reproduced by the absence of high density cold gas in the hot haloes of the most massive (and most biased) galaxies."," In doing so, they demonstrate that the observed anti-correlation of and bias at z=0.6 may be reproduced by the absence of high density cold gas in the hot haloes of the most massive (and most biased) galaxies."968" Tinker&Chen(2010) built on the Tinker&Chen(2008) model to incorporate the observed redshift evolution of the number density, thus enabling predictions of the absorber halo occupation distribution as a function of redshift."," \citet{TC10} built on the \citet{TC08} model to incorporate the observed redshift evolution of the number density, thus enabling predictions of the absorber halo occupation distribution as a function of redshift."969 The number density of absorbers in the SDSS has been shown to be roughly constant with redshift (Nestoretal.2005;Prochter2006a;Lundgren 2009).," The number density of absorbers in the SDSS has been shown to be roughly constant with redshift \citep{Nestor05,Prochter06,L09}."970". Reconciling this non-evolving number density within the context of hierarchical growth, which produces fewer haloes at a fixed mass at higher redshifts, requires evolution in the distribution of cold gas in haloes."," Reconciling this non-evolving number density within the context of hierarchical growth, which produces fewer haloes at a fixed mass at higher redshifts, requires evolution in the distribution of cold gas in haloes."971 One can achieve this effect by varying either the effective gas radius of the haloes or the typical absorber halo mass., One can achieve this effect by varying either the effective gas radius of the haloes or the typical absorber halo mass.972" As detailed in Tinker&Chen (2010), the observational outcomes of these two scenarios are degenerate in number density, but they diverge in predicted bias over a range in redshift."," As detailed in \citet{TC10}, the observational outcomes of these two scenarios are degenerate in number density, but they diverge in predicted bias over a range in redshift."973 In Figure 5 we overlay theMgII bias measurement from this work onto the curves of projected bias evolution calculated separately for models of evolving gas radius and halo mass from Tinker&Chen(2010)., In Figure 5 we overlay the bias measurement from this work onto the curves of projected bias evolution calculated separately for models of evolving gas radius and halo mass from \citet{TC10}.974" The bias measurement at z~1 suggests a preference for the model of gas radius evolution, though the error on our measurement is still too large to rule out the mass evolution model."," The bias measurement at $\sim$ 1 suggests a preference for the model of gas radius evolution, though the error on our measurement is still too large to rule out the mass evolution model."975"It is importantto note that the model curves of Tinker&Chen have been calculated for an absorber sample with W,»1À , whereas our data extends to a lower equivalent","It is importantto note that the model curves of \citet{TC10} have been calculated for an absorber sample with $_{r}>$ , whereas our data extends to a lower equivalent"976differences in its statistical quality. may be applied to deseribe the data.,"differences in its statistical quality, may be applied to describe the data."977 All errors derived in spectral fitting are statistical errors given by their Additional uncertainties may arise from errors in the atomic data and instrumental calibration. which are not explicitly accounted for.," All errors derived in spectral fitting are statistical errors given by their Additional uncertainties may arise from errors in the atomic data and instrumental calibration, which are not explicitly accounted for."978 We investigated the X-ray properties of Altair by analyzing its light curves. spectra and emission lines and present our results obtained from different analysis methods in the respective physical context.," We investigated the X-ray properties of Altair by analyzing its light curves, spectra and emission lines and present our results obtained from different analysis methods in the respective physical context."979 To search for short-term vartability on timescales of hours to days. we investigated the temporal behavior of Altair's X-ray brightness during our 2x2 exposures. that are separated by 14 days and cover in total roughly hh of observation time.," To search for short-term variability on timescales of hours to days, we investigated the temporal behavior of Altair's X-ray brightness during our 2x2 exposures, that are separated by 14 days and cover in total roughly h of observation time."980 In refle we show the count rate in the 0.222.0 keV band as measured by the EPIC. i.e. summed PN and MOS. instrument in time-steps of half an hour.," In \\ref{lc} we show the count rate in the 2.0 keV band as measured by the EPIC, i.e. summed PN and MOS, instrument in time-steps of half an hour."981 While clearly no strong flares are present. significant variability of Altair’s X-ray brightness at a level of roughly in all individual exposures and over the total observation time.," While clearly no strong flares are present, significant variability of Altair's X-ray brightness at a level of roughly in all individual exposures and over the total observation time."982 This variability could in. principle be due to rotational modulation or caused by intrinsic variability of the X-ray emitting features: e.g. microflaring or emergence/decay of weakly active regions., This variability could in principle be due to rotational modulation or caused by intrinsic variability of the X-ray emitting features; e.g. microflaring or emergence/decay of weakly active regions.983 These possible scenarios can in principle be distinguished by studying the periodicity of the light curve or the spectral changes related to the changes in brightness., These possible scenarios can in principle be distinguished by studying the periodicity of the light curve or the spectral changes related to the changes in X-ray brightness.984 Due to Altair’s inclination of about ..only those features being sufficiently close to the surface and being located at equatorial up to intermediate latitudes would induce a rotational modulation: in contrast. emission from high and polar latitudes or very extended regions would be always visible.," Due to Altair's inclination of about ,only those features being sufficiently close to the surface and being located at equatorial up to intermediate latitudes would induce a rotational modulation; in contrast, emission from high and polar latitudes or very extended regions would be always visible."985 Spectral variations are expected to be rather minor in the case of rotational modulation and for the emergence/decay of weakly active regions. in contrast flaring should be accompanied by a spectral hardening.," Spectral variations are expected to be rather minor in the case of rotational modulation and for the emergence/decay of weakly active regions, in contrast flaring should be accompanied by a spectral hardening."986 Rotational modulation can be studied with our X-ray data that covers in total three to four stellar rotations., Rotational modulation can be studied with our X-ray data that covers in total three to four stellar rotations.987 However. the uncertainties in the relevant stellar parameters (Vsini.ἐς R) allow a range of rotation periods with values around 9.5+1 h. Further. our data is not continuous. and especially the data gap of two weeks between both observations leads to large phase uncertainties.," However, the uncertainties in the relevant stellar parameters $Vsini, i, R$ ) allow a range of rotation periods with values around $9.5\pm 1$ h. Further, our data is not continuous, and especially the data gap of two weeks between both observations leads to large phase uncertainties."988 Additionally. some restructuring of X-ray emitting surface features might have occurred over times of several days.," Additionally, some restructuring of X-ray emitting surface features might have occurred over times of several days."989 To investigate periodicity of the obtained X-ray light curves. that might be related to rotational modulation. we tested periods in the range of 113 h for both observations separately. resulting in à minimum of ten time bin pairs per period and observation.," To investigate periodicity of the obtained X-ray light curves, that might be related to rotational modulation, we tested periods in the range of 13 h for both observations separately, resulting in a minimum of ten time bin pairs per period and observation."990 We then calculated the variation of the folded light curve. weighted with the mean deviation. for each period.," We then calculated the variation of the folded light curve, weighted with the mean deviation, for each period."991" As shown in the upper panel of refpha,r..amininuan. particularly pronounced forthe firstobservation. is This indicates. that rotational modulation is indeed present and points to a distribution of active regions that persists at least over a stellar rotation."," As shown in the upper panel of \\ref{pha_hr}, a minimum, particularly pronounced for the first observation, is obtained for periods around $10\pm 1$ h. This indicates, that rotational modulation is indeed present and points to a distribution of active regions that persists at least over a stellar rotation."992 The X-ray period is fully consistent with the range of rotation periods derived above and particularly favors values the longer periods., The X-ray period is fully consistent with the range of rotation periods derived above and particularly favors values the longer periods.993 Neglecting possible differential rotation. that might even be anti-solar. our periods correspondingly suggest values in the lower range for Vsini or in the upper range for the radius.," Neglecting possible differential rotation, that might even be anti-solar, our periods correspondingly suggest values in the lower range for $Vsini$ or in the upper range for the radius."994 A period of around IOhh provides very good self-similarity of the data obtained from the first observation within errors. ie. over roughly two rotations.," A period of around h provides very good self-similarity of the data obtained from the first observation within errors, i.e. over roughly two rotations."995 Likewise 1t deseribes the second observation best. however here the scatter is much larger and the X-ray light curve appears rather irregular.," Likewise it describes the second observation best, however here the scatter is much larger and the X-ray light curve appears rather irregular."996 We suspect that short-term variability. emergence and decay of quiescent coronal features or coronal restructuring are very likely responsible for this behavior.," We suspect that short-term variability, emergence and decay of quiescent coronal features or coronal restructuring are very likely responsible for this behavior."997 The here derived X-ray period ts in the same range. but slightly larger than a low-frequency period (P 99.3hh) found in the pulsation study of ? and the periods derived from modeling of interferometric data. for example P 2299.3hh (?) and P==88.9hh (?).," The here derived X-ray period is in the same range, but slightly larger than a low-frequency period $P$ h) found in the pulsation study of \cite{buz05} and the periods derived from modeling of interferometric data, for example $P$ h \citep{sou05} and $P$ h \citep{pet06}."998" If the differences have to be attributed to systematic and measurement errors or indicate different ""physical periods remains uncertain. yet the derived periods of Altair are close to each other. supporting that rotational modulated X-ray emission is indeed present during the observation."," If the differences have to be attributed to systematic and measurement errors or indicate different `physical' periods remains uncertain, yet the derived periods of Altair are close to each other, supporting that rotational modulated X-ray emission is indeed present during the observation."999 The absence of larger or even moderate flares. which are commonly observed in more active stars. may be a chance effect. but indicates the rareness of such events In stars with shallow convection zones.," The absence of larger or even moderate flares, which are commonly observed in more active stars, may be a chance effect, but indicates the rareness of such events in stars with shallow convection zones."1000 The overall X-ray brightness remained fairly constant (mean net count rate of 0.23 cts/s vs. 0.2] ets/s) between both observations. Le. over two weeks. corresponding to more than 30 stellar rotations.," The overall X-ray brightness remained fairly constant (mean net count rate of 0.23 cts/s vs. 0.21 cts/s) between both observations, i.e. over two weeks, corresponding to more than 30 stellar rotations."1001 Further. the brightness only varies by about These findings indicate. that the. X-ray emitting features are," Further, the X-ray brightness only varies by about These findings indicate, that the X-ray emitting features are"1002increasing dependence on micro turbulence in cooler stars.,increasing dependence on micro turbulence in cooler stars.1003 An increase in temperature causes an increase 1n surface gravity and our values are 0.15 dex higher compared to the literature values., An increase in temperature causes an increase in surface gravity and our values are 0.15 dex higher compared to the literature values.1004 The comparison between the mean metallicity of our total sample of giant stars and giant stars with announced companions reveals that the companion hosting stars have a 0.13 + 0.03 dex higher metallicity than the mean metallicity of our total sample., The comparison between the mean metallicity of our total sample of giant stars and giant stars with announced companions reveals that the companion hosting stars have a 0.13 $\pm$ 0.03 dex higher metallicity than the mean metallicity of our total sample.1005 This is in agreement with the enhanced metallicity of companion hosting dwarf stars. but is based or low number statistics.," This is in agreement with the enhanced metallicity of companion hosting dwarf stars, but is based on low number statistics."1006 Rotational velocities are determined using the method described by ?.., Rotational velocities are determined using the method described by \citet{fekel1997}.1007 Stars in common between our sample and that observed by ? are used to convert FWHM of moderate lines ATI to total line broadening |kms! J., Stars in common between our sample and that observed by \citet{gray1989} are used to convert FWHM of moderate lines ] to total line broadening $^{-1}$ ].1008 We used a log g vs. Tar correlation to determine the luminosity class of the stars., We used a $\log$ g vs. $_{\rm{eff}}$ correlation to determine the luminosity class of the stars.1009 This lummosity class was subsequently used to calculate the macro turbulence. which has a different relation with temperature for different classes.," This luminosity class was subsequently used to calculate the macro turbulence, which has a different relation with temperature for different classes."1010 Our data are in agreement with those obtained by ?.. but are on average larger than the values obtained by ?..," Our data are in agreement with those obtained by \citet{gray1989}, but are on average larger than the values obtained by \citet{demedeiros1999}."1011 This is due to the different diagnosties used to determine esinf. , This is due to the different diagnostics used to determine $\varv \sin i$ 1012"In these Appendices, a number of instructional cases are considered for the polarized line profiles shapes from a Keplerian disk when the illuminating star is treated as a point source.","In these Appendices, a number of instructional cases are considered for the polarized line profiles shapes from a Keplerian disk when the illuminating star is treated as a point source."1013 This means that both stellar occultation and the finite star depolarization factor are ignored., This means that both stellar occultation and the finite star depolarization factor are ignored.1014 A consequence of this approximation is that a non-zero pprofile canonly result from the Hanle effect., A consequence of this approximation is that a non-zero profile can result from the Hanle effect.1015" Before considering polarized line profiles, Stokes-I profile shapes are derived for the case of isotropic scattering."," Before considering polarized line profiles, Stokes-I profile shapes are derived for the case of isotropic scattering."1016 These solutions form the base emissivity function from which the polarized lines are constructed., These solutions form the base emissivity function from which the polarized lines are constructed.1017" Isotropic scattering corresponds to E,=0, and it means there is no polarization from resonance line scattering."," Isotropic scattering corresponds to $E_1=0$, and it means there is no polarization from resonance line scattering."1018" Of course, that also means there is no Hanle effect, regardless of the field strength."," Of course, that also means there is no Hanle effect, regardless of the field strength."1019" Even though there is no Hanle effect, the isotropic case is useful to explore as a reference for the production of the Stokes-I line shape."," Even though there is no Hanle effect, the isotropic case is useful to explore as a reference for the production of the Stokes-I line shape."1020 The integrand for the line emission as a function of velocity shift in the observed line represents the contributions by the disk density and the Sobolev effect for the profile shape., The integrand for the line emission as a function of velocity shift in the observed line represents the contributions by the disk density and the Sobolev effect for the profile shape.1021 Allowing for ΕΙ#0 and the Hanle effect simply represents new weighting functions for non-isotropic scattering that multiply the integrand from the isotropic case., Allowing for $E_1 \neq 0$ and the Hanle effect simply represents new weighting functions for non-isotropic scattering that multiply the integrand from the isotropic case.1022 The flux of line emission at normalizedDoppler shift t; is where the factor of 2 arises from the back-front symmetry of the integration along the isovelocity zone., The flux of line emission at normalizedDoppler shift $\wz$ is where the factor of 2 arises from the back-front symmetry of the integration along the isovelocity zone.1023" As a reminder, t=σ΄} and to=w?."," As a reminder, $t=\varpi^{-1}$ and $t_0=\wzz$."1024 Again the preceding expression is only valid in the point star approximation., Again the preceding expression is only valid in the point star approximation.1025 The power law exponent m is from the surface density distribution that is assumed to be a power law of the form X=ρα.," The power law exponent $m$ is from the surface density distribution that is assumed to be a power law of the form $\Sigma = \Sigma_0\, \varpi^{\rm -m}$."1026 This formulation leads to symmetric double-peaked line profile shapes for m> 2., This formulation leads to symmetric double-peaked line profile shapes for $m>2$ .1027 Larger values of m result in, Larger values of $m$ result in1028"As another comparison. the last four columns in Table 1 give analytically computed results of w, and A, for Riemann $-tvpe ellipsoids having the specified axes ratios b/a and e/o. as well as the ratio between numerical and analytical results (only models with a«5be have been computed).","As another comparison, the last four columns in Table \ref{n0} give analytically computed results of $\omega_{a}$ and $\lambda_{a}$ for Riemann S-type ellipsoids having the specified axes ratios $b/a$ and $c/a$, as well as the ratio between numerical and analytical results (only models with $a < b < c$ have been computed)."1029 To obtain these analvlical results. we used standard incomplete elliptic integrals to evaluate (he potential field of a homogeneous ellipsoid 1987.seetheirTable 2-2)..," To obtain these analytical results, we used standard incomplete elliptic integrals to evaluate the potential field of a homogeneous ellipsoid \cite[see their Table 2-2]{BT87}. ."1030 In virtually all cases. the error in w is < a few percent.," In virtually all cases, the error in $\omega$ is $\lesssim$ a few percent."1031 On the other hand. for the direct configurations the fractional error in A can be as large as5056 the error is largest [or models whose A almost vanishes.," On the other hand, for the direct configurations the fractional error in $\lambda$ can be as large as; the error is largest for models whose $\lambda$ almost vanishes."1032 This is understandable because a small variation in (he magnitude of A will result in a large percentage cdillerence when A is approaching zero., This is understandable because a small variation in the magnitude of $\lambda$ will result in a large percentage difference when $\lambda$ is approaching zero.1033 Also some mismatch between numerical and analvtical results may arise because the discrete nature of our grids makes it difficult [or us to define an exact axis ratio for our moclels., Also some mismatch between numerical and analytical results may arise because the discrete nature of our grids makes it difficult for us to define an exact axis ratio for our models.1034 For compressible equations of state. we can compare our results with those presented bv LaiRasio&Shapiro(1993).. who used an energy variational method (ο. determine approximate equilibrium configurations for ellipsoidal sell-gravitating svstems.," For compressible equations of state, we can compare our results with those presented by \cite{LRS93}, who used an energy variational method to determine approximate equilibrium configurations for ellipsoidal self-gravitating systems."1035 Lai.Rasio&Shapiro(1993) assumed that the iso-densitv contours of each configuration are concentric ellipses. which as we have already. emphasized is not the case.," \cite{LRS93} assumed that the iso-density contours of each configuration are self-similar concentric ellipses, which as we have already emphasized is not the case."1036 In (heir Table 4. for compressible analogues of Jacobi ellipsoids. a configuration with b/a=0.75 and cfa=0.4983 has T'/|W|=0.1407: the closest match among our 7=0.5 models (see Table 2)) with b/a=0.74 and c/a=0.487 has 177Wy=0.135.," In their Table 4, for compressible analogues of Jacobi ellipsoids, a configuration with $b/a=0.75$ and $c/a=0.4983$ has $T/|W|=0.1407$; the closest match among our $n=0.5$ models (see Table \ref{n05}) ) with $b/a=0.74$ and $c/a=0.487$ has $T/|W|=0.135$."1037 Other quantities also match well. alter a proper normalization is applied as prescribed by their equation (3.27).," Other quantities also match well, after a proper normalization is applied as prescribed by their equation (3.27)."1038B We omit more comparisons here and leave the rest to interested readers., We omit more comparisons here and leave the rest to interested readers.1039 Other models (hat are of interest are the /——2 inrotational sequence 1993)., Other models that are of interest are the $f=-2$ irrotational sequence \citep{UE98}.1040. It turns out that the / parameter is very sensitive to (he axis ratios. so we could not build models that are very close to the irrotational sequence due to the discrete nature of our computational erids.," It turns out that the $f$ parameter is very sensitive to the axis ratios, so we could not build models that are very close to the irrotational sequence due to the discrete nature of our computational grids."1041 Our model that comes closest has f.=—1.9: it is a direct configuration with n=0.5 ancl its other parameters are listed in Table 2.., Our model that comes closest has $f=-1.9$; it is a direct configuration with $n=0.5$ and its other parameters are listed in Table \ref{n05}.1042 We have presented a new method to construct 3D models for Riemann S-tvpe ellipsoids (n= 0). and the method has been extended (o construct compressible counterparts (hat share the same velocity. field as that of Riemann S-tvpe ellipsoids.," We have presented a new method to construct 3D models for Riemann S-type ellipsoids $n=0$ ), and the method has been extended to construct compressible counterparts that share the same velocity field as that of Riemann S-type ellipsoids."1043 With this method. we have been able to build 3D models (hat cover almost (he entire parameterspace of Riemann 5-Lvpe ellipsoids.," With this method, we have been able to build 3D models that cover almost the entire parameterspace of Riemann S-type ellipsoids."1044 Our results are in good agreement will previous studies., Our results are in good agreement with previous studies.1045 We expect (hat this method can also be straightforwardly extended to the construction of Roche-Riemann, We expect that this method can also be straightforwardly extended to the construction of Roche-Riemann10462007).,.1047. This must be related to the decreasing star formation activity over the same period (e.g..LeFlochetal. 2005).. aud the production of such quiesceut ealaxies through the truncation of star formation (6.9..Faberetal.2007:Bellct 2007): the color scatter among quiescent galaxies aud its evolution are in precise agreement with such a scenario (Ruhlancdetal.2009).," This must be related to the decreasing star formation activity over the same period \citep[e.g.,][]{lefloch05}, and the production of such quiescent galaxies through the truncation of star formation \citep[e.g.,][]{faber07, bell07}; the color scatter among quiescent galaxies and its evolution are in precise agreement with such a scenario \citep{ruhland09}."1048. There ares however. quiesceut galaxies at all redshifts 2S13 that are more massive than the most massive star-forming ealaxics.," There are, however, quiescent galaxies at all redshifts $z\lesssim 1.3$ that are more massive than the most massive star-forming galaxies."1049 This imiplics that star formation in the most massive galaxies was truucated even earlier. and/or that merecrs play an inportaut role in producing assive ealaxies.," This implies that star formation in the most massive galaxies was truncated even earlier, and/or that mergers play an important role in producing massive galaxies."1050 Evidence for the carly formation of massive galaxies is provided by their old stellar populations., Evidence for the early formation of massive galaxies is provided by their old stellar populations.1051" Ποπονα, we need to bear in πιά that there can be a large difference between the age of the stellaz population aud the assembly age. especially if mergers are miportant. as is the case in a hierarchical framework for galaxy formation (DeLuciaetal.2007)."," However, we need to bear in mind that there can be a large difference between the age of the stellar population and the assembly age, especially if mergers are important, as is the case in a hierarchical framework for galaxy formation \citep{delucia07}."1052. Hence. the muuber density evolution of galaxies is iniportant iu constraining their assembly history.," Hence, the number density evolution of galaxies is important in constraining their assembly history."1053 Measure this is cifficult because of its seusitivity to the luminosity evolution correction. especially for massive galaxies at the exponential cut-off of the mass function.," Measuring this is difficult because of its sensitivity to the luminosity evolution correction, especially for massive galaxies at the exponential cut-off of the mass function."1054 As a result. there is no conseusus among the currently available measurements (Cimattietal. 2008).," As a result, there is no consensus among the currently available measurements \citep{cimatti06, wake06, brown07,1055 cool08}."1056. Given these difficulties. other observations have beeu used. to either directly or inclirectly coustrain the assembly of galaxies.," Given these difficulties, other observations have been used to either directly or indirectly constrain the assembly of galaxies."1057 Merging activity amone the massive ealaxv population is observed (c.e..vanDoldctal. 2008)... and has been shown to produce a color-magnitude relation that is iu agreement with observations (Skeltouetal. 2009).," Merging activity among the massive galaxy population is observed \citep[e.g.,][]{vandokkum99, vandokkum05, bell06a, bell06b, lin08}, and has been shown to produce a color-magnitude relation that is in agreement with observations \citep{skelton09}."1058 Tlowever. its cosnological relevance has always been difficult to determine. given the nucertaimtics iu converting observed iuereer fractions to merger rates and the associated growth in mass.," However, its cosmological relevance has always been difficult to determine, given the uncertainties in converting observed merger fractions to merger rates and the associated growth in mass."1059 An independent and indirect indication that massive ealaxies undergo continuous evolution is provided by the recent result that ligh-+vedshift quicscent sealaxies are substantially smaller than local galaxies with the sale nns (see.vanderWeletal.2008.audrefer-ences therein).," An independent and indirect indication that massive galaxies undergo continuous evolution is provided by the recent result that high-redshift quiescent galaxies are substantially smaller than local galaxies with the same mass \citep[see,][and references therein]{vanderwel08c}."1060 This strongly suggests that mergers are nmuportant (sec.e.g.vanderWeletal.2009).. and that the assembly of massive galaxies is contiuuius up uutil the preseut dav.," This strongly suggests that mergers are important \citep[see,1061e.g.,][]{vanderwel09a}, and that the assembly of massive galaxies is continuing up until the present day."1062 Another indirect. vet powerful. constraint is provided by the evolution iu the clustering and halo occupation distribution of red galaxies (Whiteetal.2007:ConroyBrown 2008):: the evolution in the clustering strength of red salaxies is slower than expected in the absence of mereine.," Another indirect, yet powerful, constraint is provided by the evolution in the clustering and halo occupation distribution of red galaxies \citep{white07, conroy07,1063 brown08}: the evolution in the clustering strength of red galaxies is slower than expected in the absence of merging."1064 Iu this Letter we address the question whether major moreie is the dominant mechauisu for the production of verv massive. quicscent galaxies.," In this Letter we address the question whether major merging is the dominant mechanism for the production of very massive, quiescent galaxies."1065 The aremuent that we invoke is simply that major mereie ecucrally leads to roundoer galaxies., The argument that we invoke is simply that major merging generally leads to rounder galaxies.1066" Au analysis ofthe shape distribution of quiescent ealaxies can therefore coustrain the importance of ποσο,", An analysis of the shape distribution of quiescent galaxies can therefore constrain the importance of merging.1067 Since niereiug among galaxies with mass ratios of =3 is the onlv known mechanisia to produce round galaxies (see Section 3 for further discussiou). this is a powerful test.," Since merging among galaxies with mass ratios of $\lesssim 3$ is the only known mechanism to produce round galaxies (see Section 3 for further discussion), this is a powerful test."1068 The disadvantage of this method. compared to those mentioned above. is that no information about the time scale and epoch of galaxy assembly can be ferred.," The disadvantage of this method, compared to those mentioned above, is that no information about the time scale and epoch of galaxy assembly can be inferred."1069 Vincent&Ryden(2005) and Paclilla&Strauss(2008) were the first to systematically study the axial ratio distribution. p(b/a). of a laree ΙΟ of galaxies. selected fron the Sloan Digital Sky Survey (SDSS).," \citet{vincent05} and \citet{padilla08} were the first to systematically study the axial ratio distribution, $p(b/a)$, of a large number of galaxies, selected from the Sloan Digital Sky Survey (SDSS)."1070 Through a detailed analysis. they inter the intrinsic shape distribution and the effect of extinction.," Through a detailed analysis, they infer the intrinsic shape distribution and the effect of extinction."1071" Both divide the sample iuto ""elliptical aud ‘spiral’ galaxies. and coufimed that luminous elliptical galaxies arc. on average. rounder and triaxial. compared to low-huninosity “ellipticals”. which are more elongated ancl oblate (Daviesctal.1983:Frauxet1991).. and display disky isophotes (Jorgensen&Fraux1991)."," Both divide the sample into 'elliptical' and 'spiral' galaxies, and confirmed that luminous 'elliptical' galaxies are, on average, rounder and tri-axial, compared to low-luminosity 'ellipticals', which are more elongated and oblate \citep{davies83, franx91}, and display disky isophotes \citep{jorgensen94}."1072. This phenomenon is not recent: Toldenetal.(2009) showed that this trend persists at least out to :~1.," This phenomenon is not recent: \citet{holden09a}1073 showed that this trend persists at least out to $z\sim 1$."1074 Tere we preseut a complementary. modified analysis. focusing 6i ο) as a function of stellar mass for quiescent. Le. non-star-foriumng. galaxies.," Here we present a complementary, modified analysis, focusing on $p(b/a)$ as a function of stellar mass for quiescent, i.e., non-star-forming, galaxies."1075 Because mas-to-liehlt ratios are well coustrained bv broad-baud colors for mescent galaxies. stellar iüass estimates are robust.," Because mass-to-light ratios are well constrained by broad-band colors for quiescent galaxies, stellar mass estimates are robust."1076 This is essential for our purposes. as we are interested in the most massive objects; Le.. those that populate the exponential tail of the mass function.," This is essential for our purposes, as we are interested in the most massive objects, i.e., those that populate the exponential tail of the mass function."1077 Furthermore. as opposed to previous studies. we pre-select ealaxies independent of their photometric properties.," Furthermore, as opposed to previous studies, we pre-select galaxies independent of their photometric properties."1078 Our sliape-independent. spectroscopic selection criteria circumvent the biases that are poteutially introduced by selecting ealaxies by their morphological’ properties. or sone surface brightness profile.," Our shape-independent, spectroscopic selection criteria circumvent the biases that are potentially introduced by selecting galaxies by their 'morphological' properties, or some pre-defined surface brightness profile."1079 With this sample. for which we have determined wial ratios from our own fits to two-dimensional light istributious. we address the following specific questions.," With this sample, for which we have determined axial ratios from our own fits to two-dimensional light distributions, we address the following specific questions."1080 Are hiel-ass. quiescent ealaxies rounder than low-lass quiescent galaxies?," Are high-mass, quiescent galaxies rounder than low-mass quiescent galaxies?"1081 Tf so. is there a mass Ht at which p(b/a) distinctly changes. aud above which disk-dominated are completely absent?," If so, is there a mass limit at which $p(b/a)$ distinctly changes, and above which disk-dominated are completely absent?"1082 Such evidence would nuply that the only evolutionary path to such masses is a cisk-destroving mechanisui. L6. major mereiug.," Such evidence would imply that the only evolutionary path to such masses is a disk-destroying mechanism, i.e., major merging."1083 We select a sample of 17.180 quiesceut galaxies from Data Release 6 of the SDSS (Adchnan-\leCart2008).," We select a sample of 17,480 quiescent galaxies from Data Release 6 of the SDSS \citep{adelman08}."1084. Our sample iucludes galaxies at redshifts 0.01<hy2«0.08 without detectable |OII| aud Πα enissiou lines., Our sample includes galaxies at redshifts $0.04<z<0.08$ without detectable $[\rm{OII}]$ and $\rm{H}\alpha$ emission lines.1085 The selection criteria are described aud motivated in full by Gravesetal.(20001 but as opposed to that work. we do not exclude galaxies with a low concentration iudex and galaxies that are fit better by an exponcutial profile than by a deVaucouleurs(1918). profile. because this may exclude quiesceut. vet disk-like galaxies. which are obviously relevaut for quantifvius p(b/a) of quiescenut ealaxies.," The selection criteria are described and motivated in full by \citet{graves09b}; but as opposed to that work, we do not exclude galaxies with a low concentration index and galaxies that are fit better by an exponential profile than by a \citet{devaucouleurs48}1086 profile, because this may exclude quiescent, yet disk-like galaxies, which are obviously relevant for quantifying $p(b/a)$ of quiescent galaxies."1087 As a consequence. our sample may include ealaxies with star formation im an exteuded disk outside the SDSS spectroscopic fiber.," As a consequence, our sample may include galaxies with star formation in an extended disk outside the SDSS spectroscopic fiber."1088 This effect. however. does not affect our mn conclusion that quiesceut massive ealaxies with pronmünenut disks are extremely rare (see Section ??)).," This effect, however, does not affect our main conclusion that quiescent massive galaxies with prominent disks are extremely rare (see Section \ref{res}) )."1089 Rather. such a bias works iu the opposite direction in the sense that it would lead to the mistaken inclusion of galaxies with large disks.," Rather, such a bias works in the opposite direction in the sense that it would lead to the mistaken inclusion of galaxies with large disks."1090 The exclusion of all galaxies with enmuüssiou lines. also excludes quiescent galaxies with active galactic nuclei., The exclusion of all galaxies with emission lines also excludes quiescent galaxies with active galactic nuclei.1091 Their uuuber. however. is s2all. aud make up a sinall fraction of the population (c.e..Pasqualietal.2009) ," Their number, however, is small, and make up a small fraction of the population \citep[e.g.,][]{pasquali09a} "1092"More than twenty vears of. exteusive. work onu cosinological. N-body- πα.Ίσα. simulations. have provided numerous detailed predictions for the structure of dark matter halos in the hierarchical clustering scenario,Navarro.",More than twenty years of extensive work on cosmological N-body numerical simulations have provided numerous detailed predictions for the structure of dark matter halos in the hierarchical clustering scenario.1093Freul.&White(1997.hereafterNEW). preceded by the pioneering efforts of Quinnetal.(1986):Freukrenetal.(1992):Crone(199 1).. sugeested a siuple fitting formmla to describe the spherically averaged density profile of isolated dark matter halos in virial equilibrium.,"\citet[][hereafter NFW]{NFW}1094 preceded by the pioneering efforts of \citet{Quinn86,Frenk88,DC91,Warren92,Crone94}, suggested a simple fitting formula to describe the spherically averaged density profile of isolated dark matter halos in virial equilibrium."1095 Since then nmuuerous siuulations were done for mauv relaxed halos of different masses and in different cosmologies., Since then numerous simulations were done for many relaxed halos of different masses and in different cosmologies.1096 The NEW aualvtical density profile has two paraicters: the characteristic deusitv p. auc the radius à., The NFW analytical density profile has two parameters: the characteristic density $\rho_{\rm s}$ and the radius $r_{\rm s}$.1097" Iustead of these parameters. one can use the virial mass of the halo. Mag. aud the couceutratioun. €= Πίος,"," Instead of these parameters, one can use the virial mass of the halo, $\Mvir$, and the concentration, $C\equiv\Rvir/r_{\rm s}$ ."1098 Here the mass Mags and the correspoudiug radius B are defined as the mass and the radius within which the spherically averaged overdensity is equal to some specific value., Here the mass $\Mvir$ and the corresponding radius $\Rvir$ are defined as the mass and the radius within which the spherically averaged overdensity is equal to some specific value.1099 For thestandard cosmological model with the cosmological coustant ACDM andxumanmneters O4Ü= 0.3. Q4A= 0.7. and /=0.7 we," For thestandard cosmological model with the cosmological constant $\LCDM$ andparameters $\Omega_{0}=0.3$ , $\Omega_{\Lambda} = 0.7$ , and $h = 0.7$ we"1100In Figure 4 (left panel). the hard photon index distribution in NLSv1 sample (shaded: histogram) is compared. to the DBLSy1 one (IBIS fourth catalogue. Molina et al.,"In Figure \ref{figure=foton} (left panel), the hard photon index distribution in NLSy1 sample (shaded histogram) is compared to the BLSy1 one (IBIS fourth catalogue, Molina et al."1101 in. preparation)., in preparation).1102 The hare photon index is broadly distributed in our sample ranging from  1.3 to very steep values such as  3.6. with a mean value «bou.1005> = 2340.7 (see Table 2)).," The hard photon index is broadly distributed in our sample ranging from $\sim$ 1.3 to very steep values such as $\sim$ 3.6, with a mean value $<$$\Gamma_{20-100 keV}$$>$ $=$ $\pm$ 0.7 (see Table \ref{table=hard}) )."1103 This value is consistent within errors with the one previously found. in a smaller sample of tive NLSy1 (κοςtoon = 2.60.3.Alalizia ct al.," This value is consistent within errors with the one previously found in a smaller sample of five NLSy1 $<$$\Gamma_{20-100 keV}$$>$ $=$ $\pm$ 0.3,Malizia et al."1104 2008)., 2008).1105 Phe BLSy1 mean hard photon index «VouLOUbYut = 2040.2) is only slightlv [latter and consistent with the Αν] «Lejgoopevc as dndeced a Ixolmogorov-Smirnov. (Ix-8) test (probability P= of a ranclom result) indicates. that the distributions are not significantly dilleren.," The BLSy1 mean hard photon index $<$$\Gamma_{20-100 keV}$$>$ $=$ $\pm$ 0.2) is only slightly flatter and consistent with the NLSy1 $<$$\Gamma_{20-100 keV}$$>$, as indeed a Kolmogorov-Smirnov (K-S) test (probability $P = 0.030$ of a random result) indicates that the distributions are not significantly different."1106 Phis is also in agreement with estimates. [rom tvpe 1 Sevlert spectra. from οBAT (Loulookey 5 2-28+0.11. Ajello et al.," This is also in agreement with estimates from type 1 Seyfert spectra from /BAT $<$$\Gamma_{20-100 keV}$$>$ $=$ $\pm$ 0.11, Ajello et al."1107 2008)., 2008).1108 From this analysis. the average photon index of NLSy1 at hard X-rays appears not to be steeper than in BLSw however. we should consider that faint steep spectrum sources1. may be missed by hard X-ray surveys.," From this analysis, the average photon index of NLSy1 at hard X-rays appears not to be steeper than in BLSy1, however, we should consider that faint steep spectrum sources may be missed by hard X-ray surveys."1109 ὃν fitting the 0.8-100 keV. broad-band spectrum (as in ‘Table 6 and including data from literature for NGC 4051. Ark 783 and NGC 5506). the resulting photon index varies rom ~ 1.5 to — 2.6 with a mean of «Lusi005Vv = 2.0zE0.3. consistent with the typical values found. for this class of sources (e... Leighlv 1999).," By fitting the 0.3-100 keV broad-band spectrum (as in Table \ref{table=best} and including data from literature for NGC 4051, Mrk 783 and NGC 5506), the resulting photon index varies from $\sim$ 1.5 to $\sim$ 2.6 with a mean of $<$$\Gamma_{0.3-100 keV}$$>$ $=$ $\pm$ 0.3, consistent with the typical values found for this class of sources (e.g., Leighly 1999)."1110 We should note that the 3-100 keV. spectral fit parameters are more representative of the spectra below LO keV. for the larger statistics. especially when data are used.," We should note that the 0.3-100 keV spectral fit parameters are more representative of the spectra below 10 keV, for the larger statistics, especially when data are used."1111 In a complete nud X-ray selected sample of BLSv1. Molina et al. (," In a complete hard X-ray selected sample of BLSy1, Molina et al. ("11122009) ound a Blatter mean value «Leopounce  1.3 (a = 0.3). confirming the evidence that NLSy1 tend to have steeper hoton indeces. as also confirmed in Bianchi et al. (,"2009) found a flatter mean value $<$$\Gamma_{20-100 keV}$$>$ $\sim$ 1.7 $\sigma$ $=$ 0.2), confirming the evidence that NLSy1 tend to have steeper photon indeces, as also confirmed in Bianchi et al. ("111320092). or an X-ray selected sample.,"2009a), for an X-ray selected sample."1114 Pherefore. no clear separation tween NLSy1I and BLSy average hard X-ray photon index is found. as instead: observed.1 for the broad-band photon inclex.," Therefore, no clear separation between NLSy1 and BLSy1 average hard X-ray photon index is found as instead observed for the broad-band photon index."1115 In Figure 4 (right panel) we plot the hard X-ray versus the 0.3-100. keV. broad. band X-ray photon index., In Figure \ref{figure=foton} (right panel) we plot the hard X-ray versus the 0.3-100 keV broad band X-ray photon index.1116 The one-to-one regression line is also drawn., The one-to-one regression line is also drawn.1117 The steeper hard. X-ray. photon indeces are also clear from this. plot sugeesting the possible presence of a high energy. cut-oll., The steeper hard X-ray photon indeces are also clear from this plot suggesting the possible presence of a high energy cut-off.1118 Llowever. with the present data we were able to constrain this parameter only in Swift J2127.415654 (eaapp = 49.D keV). in agreement with our previous estimate (Malizia et al.," However, with the present data we were able to constrain this parameter only in Swift J2127.4+5654 $_{cut-off}$ $=$ $^{+49}_{-17}$ keV), in agreement with our previous estimate (Malizia et al."1119 2008) and with the PLN measurement. (Miniutti et al., 2008) and with the PIN measurement (Miniutti et al.1120 2000)., 2009).1121 A very steep hard. rav photon index is measured in Ark 766 (E — 2.9) and indeed. a spectral decrease is evident at ~ 50 keV from the PIN data (Yurner ct al.," A very steep hard X-ray photon index is measured in Mrk 766 $\Gamma$ $\sim$ 2.9) and indeed, a spectral decrease is evident at $\sim$ 50 keV from the PIN data (Turner et al."1122 2007)., 2007).1123 In. AIrk 783. LRAS 15091-2107 and in Lk. 16385-2057 the hard. X-ray photon index is much steeper than the SRP and one.," In Mrk 783, IRAS 15091-2107 and in IGR J16385-2057 the hard X-ray photon index is much steeper than the XRT and one,"1124on the apparent scatter in the broad-band spectra.,on the apparent scatter in the broad-band spectra.1125 These errors are included in all tables and figures., These errors are included in all tables and figures.1126" For analysis, we grouped our data into “epochs”, with observations included in a given epoch separated by no more than three days."," For analysis, we grouped our data into “epochs”, with observations included in a given epoch separated by no more than three days."1127 Tables 1 and 2 list the flux densities at each epoch., Tables \ref{srcTab1} and \ref{srcTab2} list the flux densities at each epoch.1128" On days 40.6 and 61.4, we obtained observations of wwith the Swift satellite etal. of duration 3.0 and 5.5 ks."," On days 40.6 and 61.4, we obtained observations of with the satellite \citep{Gehrels04} of duration 3.0 and 5.5 ks."1129 Each observation(Gehrels resulted2004) in an exposure with the X-ray telescope (XRT) and an image with the Ultraviolet Optical Telescope (UVOT)., Each observation resulted in an exposure with the X-ray telescope (XRT) and an image with the Ultraviolet Optical Telescope (UVOT).1130 We used theSwift data analysis routines in HEASoft version 6.10 throughout our analysis., We used the data analysis routines in HEASoft version 6.10 throughout our analysis.1131" The UVOT images were both obtained with the UVW1 filter, which has a central wavelength of 2600 and a FWHM of 693 etal. 2008)."," The UVOT images were both obtained with the UVW1 filter, which has a central wavelength of 2600 and a FWHM of 693 \citep{Poole08}."1132. wwas not detected in either UVOT image with a limiting UVWI magnitude of 20.36 (20.65) in the first (second) observation., was not detected in either UVOT image with a limiting UVW1 magnitude of 20.36 (20.65) in the first (second) observation.1133 The XRT was operated in photon counting mode during both observations., The XRT was operated in photon counting mode during both observations.1134" We produced cleaned level 2 event files by running the XRT reduction pipeline on the level 1 event files, retaining events with grades 0-12."," We produced cleaned level 2 event files by running the XRT reduction pipeline on the level 1 event files, retaining events with grades 0–12."1135 We then used XSelect v2.4 to create spectra for each dataset., We then used XSelect v2.4 to create spectra for each dataset.1136" Source counts were extracted from 30""--radius circular region centered onAql; background counts were extracted from a 95""--radius source free region.", Source counts were extracted from -radius circular region centered on; background counts were extracted from a -radius source free region.1137" The ancilliary response files (ARF) were generated using the taskxrtmkarf, and were corrected for hot pixels and dead columns using an exposure map of each observation."," The ancilliary response files (ARF) were generated using the task, and were corrected for hot pixels and dead columns using an exposure map of each observation."1138" Finally, we used the most recent response matrix file (RMF) appropriate for PC mode and event grades 0-12 from theSwift calibration database."," Finally, we used the most recent response matrix file (RMF) appropriate for PC mode and event grades 0–12 from the calibration database."1139" Aql’ss light curves are characterized by an initial steep rise until around day 50 (S,ο.1?? at 32.1 GHz), at which point there is a turnover and rapid decay at the lower frequencies (1-9 GHz) and flattening at 19 GHz and higher "," s light curves are characterized by an initial steep rise until around day 50 $S_{\nu} \propto t^{3.3}$ at 32.1 GHz), at which point there is a turnover and rapid decay at the lower frequencies (1–9 GHz) and flattening at 19 GHz and higher (Figure \ref{lcPlot}) )."1140"In order to characterize the radio spectra, we fit (Figuresimple 1)).power-law models to each epoch's 4-37 GHz spectrum; the spectra, along with fits and residuals, are shown in Figure 2.."," In order to characterize the radio spectra, we fit simple power-law models to each epoch's 4–37 GHz spectrum; the spectra, along with fits and residuals, are shown in Figure \ref{specPlot}."1141" The fit residuals in some epochs suggest the presence of a spectral break, so we performed separate power-law fits to the lower (1—9 GHz) and higher (19-37 GHz) frequencies as well."," The fit residuals in some epochs suggest the presence of a spectral break, so we performed separate power-law fits to the lower (1--9 GHz) and higher (19–37 GHz) frequencies as well."1142" In order to determine whether any improvement in fit was statistically significant, we performed F-tests for each epoch."," In order to determine whether any improvement in fit was statistically significant, we performed F-tests for each epoch."1143" The addition of a second power-law is favored at less than 3o significance, so we use the results from the single power-law fits in our discussion ofAql."," The addition of a second power-law is favored at less than $3\sigma$ significance, so we use the results from the single power-law fits in our discussion of."1144". We also plot this spectral index a (S,ος v%) in the lower panel of Figure 1,, showing that the rise portion of the light curve is accompanied by flattening of the spectrum, which then steepens again witha the turnover in low-frequency flux."," We also plot this spectral index $\alpha$ $S_{\nu} \propto \nu^{\alpha}$ ) in the lower panel of Figure \ref{lcPlot}, showing that the rise portion of the light curve is accompanied by a flattening of the spectrum, which then steepens again with the turnover in low-frequency flux."1145 We report the results of our spectral fitting in Table 3.. TheSwift, We report the results of our spectral fitting in Table \ref{radFits}.1146 XRT detected 25 (47) events at the position of oonday 41 (61)., The XRT detected 25 (47) events at the position of onday 41 (61).1147" We estimate 2 (4) background counts, resulting in count rates of 7.6+1.6 (7.8+ 1.2) x107? c s-!."," We estimate 2 (4) background counts, resulting in count rates of $7.6\pm1.6$ $7.8\pm1.2$ ) $\times \,10^{-3}$ c $^{-1}$ ."1148" We fit the unbinned, non-background subtracted X-ray datawith absorbed power-law and"," We fit the unbinned, non-background subtracted X-ray datawith absorbed power-law and"1149By several new images of Hlt 7329 X and D obtained with HIS'T/NICMOS and. VLE/NACO with Ll ves epoch dillerence. we could reject (2Pla) the background hypothesis. that LER 7329 D would have been a non-moving background object unrelated to LU 7329 A. Hence. HV 7329 A and D form a common proper motion pair.,"By several new images of HR 7329 A and B obtained with HST/NICMOS and VLT/NACO with 11 yrs epoch difference, we could reject $\ge 21 \sigma$ ) the background hypothesis, that HR 7329 B would have been a non-moving background object unrelated to HR 7329 A. Hence, HR 7329 A and B form a common proper motion pair."1150 Phe possible detection of a small linear. change in separation (but no change in PA) is consistent with an on-skv 2D orbit of D around A. which is eccentric and/or inclined.," The possible detection of a small linear change in separation (but no change in PA) is consistent with an on-sky 2D orbit of B around A, which is eccentric and/or inclined."1151 Curvature in orbital motion as acceleration or deceleration would be a final proof for being gravitational bound. but is not vet detected. as in all other sub-stellar companions detected. by direct imaging. except PZ Tel B (Mugrauer et al.," Curvature in orbital motion as acceleration or deceleration would be a final proof for being gravitational bound, but is not yet detected, as in all other sub-stellar companions detected by direct imaging, except PZ Tel B (Mugrauer et al."1152 2010)., 2010).1153 The magnitude dillerence between HI 7329 A and D is All = 6.7520.10 mag and AWK. = 6.620.1 mag (Lable 1). with DMXE0.033 mag for Ht. 7329 A|D (281ASS). we eet Wo=11.640.1p mag for HX 7329 D: we obtain L=11.1+0.2 mage for IL D (from Table 1 with L=5.0 mag for Lh 7329 as AO-tvpe star with J—II-Ix-L-5.0 mag): from the magnitude cdillerence between LER. 7329 A and D in the LIST I110NV filter CFable 1): we get J = 12.0630.19 mae forHi 7329 DB. calibrated with the M9.5 dwarf BAI Boo21-02 from Persson ct al. (," The magnitude difference between HR 7329 A and B is $\Delta$ H = $6.75 \pm 0.10$ mag and $\Delta$ $_{\rm s}$ = $6.6 \pm 0.1$ mag (Table 1), with $5.008 \pm 0.033$ mag for HR 7329 A+B (2MASS), we get $_{\rm s}$ $11.6 \pm 0.1$ mag for HR 7329 B; we obtain $11.1 \pm 0.2$ mag for HR 7329 B (from Table 1 with L=5.0 mag for HR 7329 as A0-type star with J=H=K=L=5.0 mag); from the magnitude difference between HR 7329 A and B in the HST F110W filter (Table 1); we get J = $12.06 \pm 0.19$ mag for HR 7329 B, calibrated with the M9.5 dwarf BRI B0021-02 from Persson et al. ("11541998) and the NIC web site.,1998) and the NIC web site.1155 Those JL colors are Consistent with spectral type. MY-8 for LU 7320 D. With a bolometric correction of D.C.=3.10+0.05 mag (for M7-8. Golimowski ct al.," Those JHKL colors are consistent with spectral type M7-8 for HR 7329 B. With a bolometric correction of $_{\rm K} = 3.10 \pm 0.05$ mag (for M7-8, Golimowski et al."1156" 2004). and the distance towards 7329 A. we get a luminosity of 2.627£pn""0.087 for Lh » 1. For = 2500-2800 Ix (for nedM7-8. Golimowskl et al."," 2004), and the distance towards HR 7329 A, we get a luminosity of $\log (L_{\rm bol}/L_{\odot}) = -2.627 \pm 0.087$ for HR 7329 B. For $_{\rm eff}$ = 2500-2800 K (for M7-8, Golimowski et al."1157 2004 and Luhmann 1999 scale) at ~12 Alves. we then derive the mass of LR 7329. D from evolutionary tracks to be 20 to 50 Jup masses (Burrows ct al.," 2004 and Luhmann 1999 intermediate scale) at $\sim 12$ Myrs, we then derive the mass of HR 7329 B from evolutionary tracks to be 20 to 50 Jup masses (Burrows et al."1158 1997 Chabrier ct al.," 1997, Chabrier et al."1159 2000. Daralle et al.," 2000, Baraffe et al."1160 2002)., 2002).1161 LHence. IIR 7329 D is indeed a brown dwarf.," Hence, HR 7329 B is indeed a brown dwarf."1162" No additional companion candidates were detected. up to <9"".", No additional companion candidates were detected up to $\le 9 ^{\prime \prime}$.1163 Phe LR 7329 LID 181327 svstem is therefore a triple svstem. with two stars with debris disks (Hllt. 7329 and its wide companion LID 181327. Jackman Paresce 1993. Smith et al.," The HR 7329 / HD 181327 system is therefore a triple system with two stars with debris disks (HR 7329 and its wide companion HD 181327, Backman Paresce 1993, Smith et al."1164 2009. Schneider et al.," 2009, Schneider et al."1165 2006) plus one brown cwarf (LR M 13)., 2006) plus one brown dwarf (HR 7329 B).1166 With 3 Pie (Smith Terrile 1984. Lagrange et al.," With $\beta$ Pic (Smith Terrile 1984, Lagrange et al."11672 ων.)and PZ Vol (Smith et al., 2010) and PZ Tel (Smith et al.1168 2009. Biller ct al.," 2009, Biller et al."1169 201C). el al.," 2010, Mugrauer et al."1170 POLO) there are two more members of the ? Pic moving eroup. which have both a debris disk ancd a sub-stcllar companion. indicating quite a large fraction and motivating further searches.," 2010) there are two more members of the $\beta$ Pic moving group, which have both a debris disk and a sub-stellar companion, indicating quite a large fraction and motivating further searches."1171 We thank the ESO Paranal Team and ESO Users Support eroup., We thank the ESO Paranal Team and ESO Users Support group.1172 RN. CG. anc POBS wish to acknowledge Deutsche l'orschungsgemeinschaft (DEC) for grant. NE 515 / 30-1.," RN, CG, and TOBS wish to acknowledge Deutsche Forschungsgemeinschaft (DFG) for grant NE 515 / 30-1."1173 We used Simbad. ancl Vizier and. archival data from the ESO. LUST. and 2ALASS.," We used Simbad and Vizier and archival data from the ESO, HST, and 2MASS."1174 We thank Andreas Seifahrt for his PSI subtraction routine written in LOL., We thank Andreas Seifahrt for his PSF subtraction routine written in IDL.1175 LIST data were obtained from the cata archive at. the. Space Telescope. Institute. which is operated. by the association of Universities. for Research in Astronomy. Inc. under the NASA contract NAS 5-26555.," HST data were obtained from the data archive at the Space Telescope Institute, which is operated by the association of Universities for Research in Astronomy, Inc. under the NASA contract NAS 5-26555."1176 We would also like to thank Alexander Ixrivov and Nina Vetzlalf for valueable discussion about cebris disks and runaway stars. respectively.," We would also like to thank Alexander Krivov and Nina Tetzlaff for valueable discussion about debris disks and runaway stars, respectively."1177period reported by Schwarz et al (2004) in our data.,period reported by Schwarz et al (2004) in our data.1178 We would like to thank Gary. ΟΠΗ for use of his Fourier analvsis progranis., We would like to thank Gary Schmidt for use of his Fourier analysis programs.1179 We would also like to thank (he anonvanous referee for their many useful. comments which helped to improve this paper., We would also like to thank the anonymous referee for their many useful comments which helped to improve this paper.1180view. old galaxies with a given σ teud to be metal poor. whereas vounecr galaxies with the same σ are metal rich (c.g. Worthev 1991: Trager et al.,"view, old galaxies with a given $\sigma$ tend to be metal poor, whereas younger galaxies with the same $\sigma$ are metal rich (e.g. Worthey 1994; Trager et al."1181 20005)., 2000b).1182 Thus. recent work suggests that residuals from the color-magnitude relation are age indicators. whereas studies of chemical abundances sugeest that velocity dispersion is a colubination of both age and metallicity.," Thus, recent work suggests that residuals from the color-magnitude relation are age indicators, whereas studies of chemical abundances suggest that velocity dispersion is a combination of both age and metallicity."1183 One. of the goals of the preseut work is to see if these couclusions are consistent with ouc-another., One of the goals of the present work is to see if these conclusions are consistent with one-another.1184 We do this bv studyiug the joint distribution of color. huuiuositv aud velocity dispersion iu a suuple of 39320 early-type galaxies drawu from the Sloan Digital ον Survey (York et al.," We do this by studying the joint distribution of color, luminosity and velocity dispersion in a sample of 39320 early-type galaxies drawn from the Sloan Digital Sky Survey (York et al."1185 2000: Stoughton et al, 2000; Stoughton et al.1186 2002: Abazajian et al., 2002; Abazajian et al.1187 2003) database (hereafter SDSS)., 2003) database (hereafter SDSS).1188 Section ?? describes our sample., Section \ref{sample} describes our sample.1189 Section ον preseuts the color-magnitude aud color-7 relations. aud demonstrates that color-o aud maguitude-7 are the primary correlations.," Section \ref{cms} presents the color-magnitude and $\sigma$ relations, and demonstrates that $\sigma$ and $\sigma$ are the primary correlations."1190 A simple model is introduced which illustrates clearly what our measurements duply for the relations between age. moetalliityv aud velocity dispersion.," A simple model is introduced which illustrates clearly what our measurements imply for the relations between age, metallicity and velocity dispersion."1191 The mathematics associated with this model is in Appendix D.., The mathematics associated with this model is in Appendix \ref{algebra}.1192 Section ?? uses the model to derive a umber of consequences of the observed coloranagnitude-o relations. and summarizes our findings.," Section \ref{discuss} uses the model to derive a number of consequences of the observed $\sigma$ relations, and summarizes our findings."1193 Throughout. we assimue that {10= rülans + Mpe+ in a universe with Qy=0.3 which is spatially flat.," Throughout, we assume that $H_0=70$ km $^{-1}$ $^{-1}$ in a universe with $\Omega_0=0.3$ which is spatially flat."1194 For our analysis we used galaxies selected from the Sloan Dieital Sky Survey (SDSS) database., For our analysis we used galaxies selected from the Sloan Digital Sky Survey (SDSS) database.1195 See York et al. (, See York et al. (11962000) for a technical παν of the SDSS project: Stoughton et al. (,2000) for a technical summary of the SDSS project; Stoughton et al. (11972002) for a description of the Early Data Release: Abazajian et al. (,2002) for a description of the Early Data Release; Abazajian et al. (11982003) ct al.,2003) et al.1199 for a description of DRI. the First Data Release: Comm et al. (," for a description of DR1, the First Data Release; Gunn et al. ("12001998) for details about the camera: Fukueita et al. (,1998) for details about the camera; Fukugita et al. (

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