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To measure robust fluxes for sources close to the confusion level. we have developed a source-extraction technique. that is based on the assumption that sources detected at 250 jm with SPIRE will also be detected in deep observations with at 24 um. This assumption was suggested by the recent studies that concluded that galaxies detected in deep 24-um surveys withSpitzer produce most of the FIRB at 160 jm (Dole et 22006) and at 250. 350 and 500 uim (Marsden et 22009). | To measure robust fluxes for sources close to the confusion level, we have developed a source-extraction technique that is based on the assumption that sources detected at 250 $\mu$ m with SPIRE will also be detected in deep observations with at 24 $\mu$ m. This assumption was suggested by the recent studies that concluded that galaxies detected in deep $\mu$ m surveys with produce most of the FIRB at 160 $\mu$ m (Dole et 2006) and at 250, 350 and 500 $\mu$ m (Marsden et 2009). |
This approach reduces the effective confusion noise by resolving some of the confusing background into the sources detected at 24 ym. Our "eross-ID method’ starts from a list of the 24-um sources found in the field covered by the 250-um image. | This approach reduces the effective confusion noise by resolving some of the confusing background into the sources detected at 24 $\mu$ m. Our `cross-ID method' starts from a list of the $\mu$ m sources found in the field covered by the $\mu$ m image. |
By using à matrix inversion technique. we then find the 250-j/m fluxes at these positions that provide the best fit to the 250-j/m image (Roseboom. in prep). | By using a matrix inversion technique, we then find the $\mu$ m fluxes at these positions that provide the best fit to the $\mu$ m image (Roseboom, in prep). |
A problem with methods like this is that there can be a large number of degenerate solutions when the surface-density of sources in the input catalogue Is large. as it is with the deep jm catalogues. | A problem with methods like this is that there can be a large number of degenerate solutions when the surface-density of sources in the input catalogue is large, as it is with the deep $\mu$ m catalogues. |
We have addressed this problem by iteratively reducing the number of sources in the input catalogue in order to find the list of input sources that provides the best fit to the 250-um image. | We have addressed this problem by iteratively reducing the number of sources in the input catalogue in order to find the list of input sources that provides the best fit to the $\mu$ m image. |
Full details of the method. its validation with simulations. and à comparison of the results with other methods of source extraction are given in Roseboom et ((in prep). | Full details of the method, its validation with simulations, and a comparison of the results with other methods of source extraction are given in Roseboom et (in prep). |
Roseboom et (Cin prep) have assessed the completeness of the eross-ID method by inserting artificial sources. onto the real images and then determining the fraction of these sources that are detected by the source-extraction method. | Roseboom et (in prep) have assessed the completeness of the cross-ID method by inserting artificial sources onto the real images and then determining the fraction of these sources that are detected by the source-extraction method. |
Fig. | Fig. |
| shows the results for the three fields as a function of flux density. | 1 shows the results for the three fields as a function of flux density. |
Using these curves to choose the flux limits. we have used the cross-ID method to extract samples of sources from regions within each field for which the optical/IR data is | Using these curves to choose the flux limits, we have used the cross-ID method to extract samples of sources from regions within each field for which the optical/IR data is |
Stern.D.. Poutanen.J. Svensson.R. 1999. ApJ. 510. Wooslev.S. 1993. ApJ 405. Zhang.D. Meésszarros.?. 2002. ApJ in press Zhang.W. Wooslev. S.. AlacFadven. A. 2002. in | Stern,B., Poutanen,J. Svensson,R. 1999, ApJ, 510, Woosley,S. 1993, ApJ 405, Zhang,B. Mésszárros,P. 2002, ApJ in press Zhang,W, Woosley, S., MacFadyen, A. 2002, in |
Among the discoveries made by theSwift satellite within a few months of its launch is the observation that a traction of long duration gamma-ray bursts (GRBs) go through an carly phase of relatively slow decline in the X-ray afterglow flux that typically starts at a few minutes after the burst and lasts for about an hour (Nouscketal.2005). | Among the discoveries made by the satellite within a few months of its launch is the observation that a fraction of long duration gamma-ray bursts (GRBs) go through an early phase of relatively slow decline in the X-ray afterglow flux that typically starts at a few minutes after the burst and lasts for about an hour \citep{Nousek05}. |
. This phase is followed by a somewhat faster and more typical flux decay that satisfies the expected relation between the temporal decline index o and the spectral index 3, where F,xv7’f°, similar to what was observed before theSwift era when the monitoring of the afterglow light curves started at least several hours after the GRB. | This phase is followed by a somewhat faster and more typical flux decay that satisfies the expected relation between the temporal decline index $\alpha$ and the spectral index $\beta$, where $F_\nu\propto \nu^{-\beta} t^{-\alpha}$, similar to what was observed before the era when the monitoring of the afterglow light curves started at least several hours after the GRB. |
The spectral index does not seem to undergo any change when the light-curve transitions (at s) from a shallow decline (42) to the "regular" decline (o3). | The spectral index does not seem to undergo any change when the light-curve transitions (at $t_{\rm break,2}\sim 10^4\;$ s) from a shallow decline $\alpha_2$ ) to the “regular” decline $\alpha_3$ ). |
It has been argued convincingly by a number of authors that the more slowly declining lightcurve, like the "regular" flux decay rate that follows it, are both produced by the shock heated circum-burst medium (Nouscketal.2005:etal.2005:Zhang 2005). | It has been argued convincingly by a number of authors that the more slowly declining lightcurve, like the “regular” flux decay rate that follows it, are both produced by the shock heated circum-burst medium \citep{Nousek05,Panaitescu05,Zhang05}. |
. The shallow X-ray flux decay is widely attributed to energy injection into the aftereglow shock, which may be caused by either a long lived activity of | The shallow X-ray flux decay is widely attributed to energy injection into the afterglow shock, which may be caused by either a long lived activity of |
(he empirical linear relationship between maser disk size and black hole mass. | the empirical linear relationship between maser disk size and black hole mass. |
The capture of clouds with column densities <107cm7 results in a non-sell-eravitating disks of the correct scale and sufficient column density to allow X-ray irracliation from the central source to reproduce thin meganmaser disks. | The capture of clouds with column densities $\la 10^{23.5}\ut cm -2 $ results in a non-self-gravitating disks of the correct scale and sufficient column density to allow X-ray irradiation from the central source to reproduce thin megamaser disks. |
By contrast. the capture of a cloud with higher column density woulcl imstead create a sell-gravitating disk giving rise to rapid star formation. | By contrast, the capture of a cloud with higher column density would instead create a self-gravitating disk giving rise to rapid star formation. |
In Paper I we showed that (his can explain the recent formation of the compact disks of stars within a fraction of a parsec of the Galactic center black hole Ser À* (Paper I: Bonnell Rice 2003: Mapelli et al. | In Paper I we showed that this can explain the recent formation of the compact disks of stars within a fraction of a parsec of the Galactic center black hole Sgr A* (Paper I; Bonnell Rice 2008; Mapelli et al. |
2008; Alig et al. | 2008; Alig et al. |
2011). | 2011). |
Note that the transition column density is sensitive to (he speed of the incoming cloud and the degree of the angular momentum cancellation. though we have chosen plausible values of these parameters. | Note that the transition column density is sensitive to the speed of the incoming cloud and the degree of the angular momentum cancellation, though we have chosen plausible values of these parameters. |
This picture relies on the presence of dense clouds close to the nucleus of galaxies. | This picture relies on the presence of dense clouds close to the nucleus of galaxies. |
The circumnuclear molecular ring CChristopher et al. | The circumnuclear molecular ring Christopher et al. |
2005). ppc [rom Ser À* in our own Galaxy. and (he circumnuclear rings found on scales of several parsecs [rom the center of numerous Sevlert galaxies suggest an ample supply of material. | 2005), pc from Sgr A* in our own Galaxy, and the circumnuclear rings found on scales of several parsecs from the center of numerous Seyfert galaxies suggest an ample supply of material. |
Recent simulations sugeest that gas supply to galactic centers is controlled by angular momen(iumnm (ransler Tom one massive gas Chunp to another during gravitational encounters (Namekata Iabe 2011). | Recent simulations suggest that gas supply to galactic centers is controlled by angular momentum transfer from one massive gas clump to another during gravitational encounters (Namekata Habe 2011). |
Some of these inward-moving clouds may interact wilh their host supermassive black 1oles, | Some of these inward-moving clouds may interact with their host supermassive black holes. |
The rate of migration of molecular material is estimated to give a potential black hole interaction rate of ~LO“vr+. | The rate of migration of molecular material is estimated to give a potential black hole interaction rate of $\sim 10^{-6}\ut yr -1 $. |
Relating these estimates to the occurrence rate of meganiaser disks requires estimates of the disk lifetimes ancl the near edge-on viewing angle needed [ου naser amplification. | Relating these estimates to the occurrence rate of megamaser disks requires estimates of the disk lifetimes and the near edge-on viewing angle needed for maser amplification. |
The single-dish detection rate of H5O megamasers in Sevífert 2 galaxies and LINERs is about1556... and about of those have kinematics consistent with orbital notion on sub-parsec scales (Lo 2005). | The single-dish detection rate of $_2$ O megamasers in Seyfert 2 galaxies and LINERs is about, and about of those have kinematics consistent with orbital motion on sub-parsec scales (Lo 2005). |
The estimated maser beaming angle. LO” MMaloney. 2002) then implies that ~20% of all Sevlert 2 galaxies possess similar disks. | The estimated maser beaming angle, $\sim 10^\circ$ Maloney 2002) then implies that $\sim20$ of all Seyfert 2 galaxies possess similar disks. |
The VLBI follow-up results suggest that roughly half may be thin. Neplerian disks. | The VLBI follow-up results suggest that roughly half may be thin, Keplerian disks. |
The ordered kinenmaties of (hese disks suggest lifetimes of hundreds of orbital periods or more: (he orbital period at ppc from a 10*M. black hole is ~3x10?vr implies lifetimes in excess of 100vr. consistent with the time seale needed to warp the disk via resonant relaxation (Alexander Bregman 2009. 2011). | The ordered kinematics of these disks suggest lifetimes of hundreds of orbital periods or more; the orbital period at pc from a $10^7\,\msun$ black hole is $\sim 3 \ee 3 \u yr $ implies lifetimes in excess of $10^6\u yr $, consistent with the time scale needed to warp the disk via resonant relaxation (Alexander Bregman 2009, 2011). |
The formation of gravitationallv-unstable disks is likely (ο be as common simply because molecular clouds in (he inner regions of galaxies tend to have column densities >107em7. | The formation of gravitationally-unstable disks is likely to be as common simply because molecular clouds in the inner regions of galaxies tend to have column densities $\ga 10^{24}\ut cm -2 $. |
These transient disks may also host meganmasers (Milosavljevié Loeb 2004). so that only a fraction of megamaser AGN maw have disks that are verv close to Neplerian enabling accurate black hole mass determinations. | These transient disks may also host megamasers (Milosavljević Loeb 2004), so that only a fraction of megamaser AGN may have disks that are very close to Keplerian enabling accurate black hole mass determinations. |
Finally. we note that the partial capture of a cloud imparts an impulse to the black hole / disk svstem of approximately AMai;;c. where e~200kins is the incident cloud velocity. | Finally, we note that the partial capture of a cloud imparts an impulse to the black hole / disk system of approximately $M_\mathrm{disk} v$, where $v\sim 200\kms $ is the incident cloud velocity. |
For plausible parameters this gives recoil velocities of ~10 20kmsὃν while this is potentially detectable. the recoil will be rapidly damped by dynamical friction on the surrounding stars. | For plausible parameters this gives recoil velocities of $\sim 10$ $20 \kms$; while this is potentially detectable, the recoil will be rapidly damped by dynamical friction on the surrounding stars, |
(his gives them an observed space densitv of 573kpe. | this gives them an observed space density of $n \sim 3
\;{\rm kpc}^{-3}$. |
This is similar (to our theoretical esimale above. given the uncertainty. associated with some of our parameters. | This is similar to our theoretical estimate above, given the uncertainty associated with some of our parameters. |
The observed stars all have a metallicity very. close to our critical metallicity. so it is nol clear if (hey. ave very low metallicity Pop IH stars or if (μον were formed by our Pop 11.9 mechanism. | The observed stars all have a metallicity very close to our critical metallicity, so it is not clear if they are very low metallicity Pop II stars or if they were formed by our Pop II.5 mechanism. |
We have shown (hat. with reasonable modeling assumptions. il is possible to produce these stars from primordial gas. and (hat (heir metallicity could be due to pollution from Pop III supernovae which (riggered (heir formation (it could also be partly due to pollution from swept up gas as (μον travel through the ealactic disk). | We have shown that, with reasonable modeling assumptions, it is possible to produce these stars from primordial gas, and that their metallicity could be due to pollution from Pop III supernovae which triggered their formation (it could also be partly due to pollution from swept up gas as they travel through the galactic disk). |
Current observational searches are beginning to eive us a clearer picture of (he number of very low metallicity stars in the galactic halo. ancl we anticipate that this will continue with future observations. | Current observational searches are beginning to give us a clearer picture of the number of very low metallicity stars in the galactic halo, and we anticipate that this will continue with future observations. |
On the theoretical side. numerical simulations can in principle reduce (he uncertainties in our Pop II.5 model by determining the values of 5. μοι and possibly also of the IME from the mass spectrum ol collapsing clumps. although this is a very challenge problem to simulate accurately, | On the theoretical side, numerical simulations can in principle reduce the uncertainties in our Pop II.5 model by determining the values of $\eta$, $f_{\rm halo}$, and possibly also of the IMF from the mass spectrum of collapsing clumps, although this is a very challenging problem to simulate accurately. |
We have investigated the history of star formation in the high redshift universe. focusing on the role plaved by very massive Pop HIE stars. | We have investigated the history of star formation in the high redshift universe, focusing on the role played by very massive Pop III stars. |
We have argued that this history is shaped bv the various feedback effects exerted bv (hose stars. resulting in three distinct epochs of slar lormation. | We have argued that this history is shaped by the various feedback effects exerted by those stars, resulting in three distinct epochs of star formation. |
The first impact Pop III stus have on (heir surroundings is radiative. as their soft UV flux dissociates the molecular hydrogen in other minihalos (ILaimanetal.1997.2000). | The first impact Pop III stars have on their surroundings is radiative, as their soft UV flux dissociates the molecular hydrogen in other minihalos \citep{HaiReeLoe97,HaiAbeRee00}. |
. Because of their copious UV emission. and the small molecular fraction (10+— *) in minihalos. this photodissociation is likely to happen quickly. and fairly completely. | Because of their copious UV emission, and the small molecular fraction $\sim10^{-4}-10^{-6}$ ) in minihalos, this photodissociation is likely to happen quickly, and fairly completely. |
A inuuerical investigation of this process was performed by Alachaceketal.(2001)... who simulated this suppression of star formation in low mass halos bv. incorporating a uniform soft UV. background into a cosmological simulation. | A numerical investigation of this process was performed by \citet{MacBryAbe01}, who simulated this suppression of star formation in low mass halos by incorporating a uniform soft UV background into a cosmological simulation. |
While Chis assumption does not allow them to follow the detailed buikl-up of the UV radiation field as Pop III star formation switches on. it does give a good indication of the overall effect. | While this assumption does not allow them to follow the detailed build-up of the UV radiation field as Pop III star formation switches on, it does give a good indication of the overall effect. |
It was found that the UV raciation ean effectively suppress star formation due to molecular cooling in low mass halos. delaving their collapse from redshift z30 to z~20. | It was found that the UV radiation can effectively suppress star formation due to molecular cooling in low mass halos, delaying their collapse from redshift $z\sim 30$ to $z\sim 20$. |
It is encouraging that both analvtic and numerical calculations agree quite well. | It is encouraging that both analytic and numerical calculations agree quite well. |
In Figure 2.. we lind that once a significant UV. radiation field has been set up by 2~30. further star formation is strongly inhibited until more massive halos start to form. and the | In Figure \ref{fig2}, we find that once a significant UV radiation field has been set up by $z\sim30$, further star formation is strongly inhibited until more massive halos start to form, and the |
The consequences of A«B on the S spectrum in region A can be derived. casily from what we mentioned before. | The consequences of $\lambda<R$ on the S spectrum in region A can be derived easily from what we mentioned before. |
The observed. peak frequency £5, remains unchanged. Fa and μμ... In our generalized model with 21 we denote quantities by the symbol ὃν where q. is defined for ?=1. (for the explicit definitions of quantities in the standard moclel we refer to Panaitescu Ixumar 2000). | The observed peak frequency $\tilde{\nu_m}$ remains unchanged, _m, and _m. In our generalized model with $\delta<1$ we denote quantities by the symbol $\tilde{q}$ ', where $q$ ' is defined for $\delta=1$, (for the explicit definitions of quantities in the standard model we refer to Panaitescu Kumar 2000). |
op . ∐↥≺⊾≼∙∪⊔↓∪∖⇁↓⊔⋏∙≟↓≻≺⋅⋜↧↓⊊↓⊔↿⋖⋅⊔⊳∖↓∣⋅∖⇁↓⊳∖∫∖. ⋠∢∕nDX. where 4ne. D' and X are respectively the external density. the magnetic field and the emitting region linear dimension in the comoving frame: therefore the observed. peak Dux. δν ds smaller than £j, by a factor 9:C'ammea:n fimesl,,. where E is the bulk Lorentz factor. | The comoving peak intensity is $I'\propto \tilde{n}'B'\lambda'$, where $\tilde{n}'$, $B'$ and $\lambda'$ are respectively the external density, the magnetic field and the emitting region linear dimension in the comoving frame; therefore the observed peak flux $\tilde{F_p}$ is smaller than $F_p$ by a factor $\delta$: ) F_p, where $\Gamma$ is the bulk Lorentz factor. |
Since the magnetic field persists for a shorter time. a smaller fraction of electrons cool significantly. before the magnetic field. disappears in that region. | Since the magnetic field persists for a shorter time, a smaller fraction of electrons cool significantly, before the magnetic field disappears in that region. |
The Lorentz factor of the electrons. whose cooling time is equal to the typical timescale of the system. ~ ds therefore higher: Suterde where Y=(1|Y)/(1Y). | The Lorentz factor of the electrons, whose cooling time is equal to the typical timescale of the system, $\tilde{\gamma_c}$ is therefore higher: _c, where ${\cal Y}=(1+Y)/(1+\tilde{Y})$. |
The observed. cooling [requeney 2xD,E is then2 | The observed cooling frequency $\tilde{\nu_c}\propto B'\,\tilde{\gamma_c}^{2}\,\Gamma$ is then. |
+ The svnchrotron. self-absorption optical thickness can be approximated as [or vy«ο. where p=c when the electrons are radiative ' vy) and p=m when they are adiabatic (7, "E | The synchrotron self-absorption optical thickness can be approximated as, for $\nu<\tilde{\nu_p}$, where $p=c$ when the electrons are radiative $\tilde{\gamma_c}<\tilde{\gamma_m}$ ) and $p=m$ when they are adiabatic $\tilde{\gamma_c}>\tilde{\gamma_m}$ ). |
We will refer to the first case as fast-cooling regime and o the second. case as slow-cooling regime. | We will refer to the first case as fast-cooling regime and to the second case as slow-cooling regime. |
The absorption requencey corresponds to το)=1 and in slow cooling regime is given by where we use Eq. | The absorption frequency corresponds to $\tilde{\tau_a}(\tilde{\nu_a})=1$ and in slow cooling regime is given by _a, where we use Eq. |
1 and Eq. 6: | \ref{eq:num} and Eq. \ref{eq:tau}; |
in fast cooling regime Eqs 5. 6 and Is (Y21) vive Eqs 3.. 7 and S are calculated in the case a.«£, (see Cranot Sari 2001 for the different possible S spectra). | in fast cooling regime Eqs \ref{eq:nuc}, \ref{eq:tau} and \ref{eq:yf} ${\cal
Y}=1$ ) give _a. Eqs \ref{eq:fp}, \ref{eq:nuas} and \ref{eq:nuaf} are calculated in the case $\tilde{\nu}_a<\tilde{\nu}_p$ (see Granot Sari 2001 for the different possible S spectra). |
=Devertheless 2,z£6, for high densities NE where £ is the isotropic equivalent energy ancl fay ds the observed. time in davs. | Nevertheless $\tilde{\nu}_a
\ge\tilde{\nu}_p$ for high densities in fast cooling and in slow cooling where $E$ is the isotropic equivalent energy and $t_{day}$ is the observed time in days. |
The inferred. external. density is reasonable for 90.1 (neeJ6:.l0'cm1 and 8.5r10[ENbu2 respectively) and for 9=0.01 (n210" and nzchiS5191day55cm 3) but for lower Dy the ensity lower limits expressed. in [eq 9 and Iq 10 are perhaps too high for the afterglow environment. | The inferred external density is reasonable for $\delta=0.1$ $\tilde{n}\ge 3.6 \times10^{4}$ $^{-3}$ and $\tilde{n}\ge 8.5\times10^{6}\,t_{day}^{5/2}$ $^{-3}$ respectively) and for $\delta=0.01$ $\tilde{n}\ge 2\times 10^{6}$ $^{-3}$ and $\tilde{n}\ge 8.5\times 10^{7}\,t_{day}^{5/2}$ $^{-3}$ ) but for lower $\delta$ the density lower limits expressed in Eq \ref{nlimif} and Eq \ref{nlimis} are perhaps too high for the afterglow environment. |
For simplicityiv we focus thereafter on ο<1% but the reacler should keep in mind that these cases hold as long as I5q 9 and Eq 10 are not satisfied. | For simplicity we focus thereafter on $\tilde{\nu}_a<\tilde{\nu}_c$ but the reader should keep in mind that these cases hold as long as Eq \ref{nlimif} and Eq \ref{nlimis} are not satisfied. |
The magnetic. energy BexnlD and the Compton‘ paranmeer are higher at the beginning of the afterglow evolution and so the electrons are more likely then to be radiative. | The magnetic energy $\frac{B'^{2}}{8\pi}\propto n \Gamma^{2}$ and the Compton parameter are higher at the beginning of the afterglow evolution and so the electrons are more likely then to be radiative. |
Since the minimum injected 5,,xL decreases with time while 5. increases as the ellicieney of raciative loss decreases. the electrons undergo a transition to the adiabatic regime. | Since the minimum injected $\gamma_m\propto \Gamma$ decreases with time while $\gamma_c$ increases as the efficiency of radiative loss decreases, the electrons undergo a transition to the adiabatic regime. |
The transition time between the fast cooling and the slow cooling reginie is Tfsos for the slow-cooling regime dominates the fireball evolution during times when the afterglow is observec (21 day) | The transition time between the fast cooling and the slow cooling regime is ^2; for the slow-cooling regime dominates the fireball evolution during times when the afterglow is observed $\gsim 1$ day). |
The accelerated: electrons that upscatter svnchrotron ((v)phelons give rise to a SSC component in the spectrum (for the SSC in the standard model see Sari Esin 2000). | The accelerated electrons that upscatter synchrotron photons give rise to a SSC component in the spectrum (for the SSC in the standard model see Sari Esin 2000). |
During the alterglow phase the medium has a TloMson optical depth of the order of 7210ὃ for p=I P? (ep is theThomsoncross section). therefore only a minor fraction of the emission. isupscattered. and. the synchrotron. [ux is not significantly altered. | During the afterglow phase the medium has a Thomson optical depth of the order of $\tilde{\tau}\simeq 10^{-5}\times \delta$ for $\tilde{n}=10^{2}$ $^{-3}$ $\sigma_T$ is theThomsoncross section), therefore only a minor fraction of the emission isupscattered, and the synchrotron flux is not significantly altered. |
Nevertheless if Y2l the electron cooling is Compton dominated and SSC /upscatters | Nevertheless if $
\tilde{Y}\ge 1$ the electron cooling is Compton dominated and SSC upscatters |
?7..) 2.. 7. ? 7, 2.. 2)). uun | \citealt{schneider04}. \citealt{vandenberk04}, \citealt{devries05}, \citealt{wilhite08}, \citealt{bauer09a}, \citealt{kelly09}, \citealt{macleod10}, \citealt{meusinger11}) |
mber of properties of the objects: c.g.. time scale of the variability. wavelength of observation. redshift. nass. and huninosity of the svstenis. | number of properties of the objects: e.g., time scale of the variability, wavelength of observation, redshift, mass, and luminosity of the systems. |
Iu particular. a strong auti-correlation has repeatedly been observed between variabilitv amplitude aud quasar Inmninosity. particularly when comparing quasars of simular mass. | In particular, a strong anti-correlation has repeatedly been observed between variability amplitude and quasar luminosity, particularly when comparing quasars of similar mass. |
Given a measurement of ai quasars variability auplitude, one can cstimate its luninosity using this clupirical relation. | Given a measurement of a quasar's variability amplitude, one can estimate its luminosity using this empirical relation. |
If the quasar were eravitationallv ensed by intervene lass. the maeuification effect would alter the observed luminosity. | If the quasar were gravitationally lensed by intervening mass, the magnification effect would alter the observed luminosity. |
Uowever. the ractional variability would remain uuaffected. as the uaeuification. which is nuiultiplicative on the huninosity. cancels. | However, the fractional variability would remain unaffected, as the magnification, which is multiplicative on the luminosity, cancels. |
Therefore. magnification due to gravitational ic shift quasars position in variabilitycsiunduosity villspace. | Therefore, magnification due to gravitational lensing will shift the quasar's position in variability-luminosity space. |
Quautificathejon of this shift constitutes Ineasurenien ο... lousing maenification. | Quantification of this shift constitutes a measurement of the quasar's lensing magnification. |
aBecause the vanabilitv-lhunminosityοἱthe relation has a large variance, the esti a edl n" Is nof precise, | Because the variability-luminosity relation has a large intrinsic variance, the estimate of a single quasar's magnification is not precise. |
mareHowever, Suglebeeanse the relation is magmineafionwell-detcrmuineda im the mean. au ofcuscuablethe of eeasarsars CADcani VICIvield‘old a sigΙΟΗΙΟsienifΠΕü TUCASTIECTACTLeasurementt of the |...leusileusine effect. | However, because the relation is well-determined in the mean, an ensemble of quasars can yield a significant measurement of the lensing effect. |
Tn this paper we take advantage of the variabilitv-luninositv relation seen iu tvpe I quasars to measure theiri maenificationqu6i due to galaxy: clusters alous: the line of seht. | In this paper we take advantage of the variability-luminosity relation seen in type I quasars to measure their magnification due to galaxy clusters along the line of sight. |
We use the Palomar-QUEST Variability Survey to measure the leusiug magnifications of 3573 quasars, aud compare the signal to that expected. from eravitational.: leusineνο⋅⋡ due to the 13.823 ↽∙⋡&alaxy clusters Eiu the Sloan; Digitalο τς,Sky Survey. ΑμαP. catalogoq (?)). | We use the Palomar-QUEST Variability Survey to measure the lensing magnifications of 3573 quasars, and compare the signal to that expected from gravitational lensing due to the 13,823 galaxy clusters in the Sloan Digital Sky Survey MaxBCG catalog\citealt{maxbcg}) ). |
We aeasure an average cluster profile shape that is | We measure an average cluster profile shape that is |
may be controlled by the aceretion and. feedback. processes occurring in the vicinity of the black hole itself. | may be controlled by the accretion and feedback processes occurring in the vicinity of the black hole itself. |
In this work. we use the combination of theSimulalion and semi-analvtic models of galaxy formation to study the fraction of galaxies that have undergone major mergers as a function of mass and cosmic epoch. | In this work, we use the combination of the and semi-analytic models of galaxy formation to study the fraction of galaxies that have undergone major mergers as a function of mass and cosmic epoch. |
We investigate whether this can be related to the demographics of black holes in the local Universe and. to the apparent disappearance of the most luminous quasar activity. in massive galaxies at late tinies. | We investigate whether this can be related to the demographics of black holes in the local Universe and to the apparent disappearance of the most luminous quasar activity in massive galaxies at late times. |
1n Sec. 2.. | In Sec. \ref{sec:simulation}, |
we briellv introduce the simulation we use and explain how galaxy mergers are tracked in the simulation. | we briefly introduce the simulation we use and explain how galaxy mergers are tracked in the simulation. |
In Sec. 3.. | In Sec. \ref{sec:merger}, |
we show that if we assume that black holes only form when galaxies undergo major merging events. then most present-day low mass galaxies are predicted not not to contain black holes and hence will not host AGN. | we show that if we assume that black holes only form when galaxies undergo major merging events, then most present-day low mass galaxies are predicted not not to contain black holes and hence will not host AGN. |
In Sec. 4.. | In Sec. \ref{sec:firstmerger}, |
we use the simulations to prediet when galaxies of cillerent masses have underdone their first major merger. | we use the simulations to predict when galaxies of different masses have underdone their first major merger. |
Conclusions and discussions are presented in the final section. | Conclusions and discussions are presented in the final section. |
The citepspringel2005. is used in this work to study the merging histories of clark matter haloes. | The \\citep{springel2005} is used in this work to study the merging histories of dark matter haloes. |
The mergingὃνo histories. of ealaxies can be inferred when the simulation is combined with semi-analvtic models that follow gas cooling. star formation. supernova and AGN feedback anc other physical processes that regulate how the baryons condense into galaxies. | The merging histories of galaxies can be inferred when the simulation is combined with semi-analytic models that follow gas cooling, star formation, supernova and AGN feedback and other physical processes that regulate how the baryons condense into galaxies. |
The follows CN.=2160* particles of mass 8.610hΝΕ. from redshift 2=127 to the present day. within a comoving box of 5005. Ape on a side. | The follows $N= 2160^3$ particles of mass $8.6\times10^{8}\,h^{-1}{\rm M}_{\odot}$ from redshift $z=127$ to the present day, within a comoving box of $500\, h^{-1}$ Mpc on a side. |
The cosmological parameters values in the simulation are consistent with the determinations from a combined analysis of the 2dEXGIS(7) and first vear WAILAP data (7)... | The cosmological parameters values in the simulation are consistent with the determinations from a combined analysis of the \citep{colless01} and first year WMAP data \citep{spergel03}. . |
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