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Third. we implemented a phenomenological model for ealactic winds in order to study the effect of outllows on DLAs. galaxies. and the intergalactic medium. (16M).
Third, we implemented a phenomenological model for galactic winds in order to study the effect of outflows on DLAs, galaxies, and the intergalactic medium (IGM).
In this model. gas particles are stochastically ciriven out ( “the dense star-forming medium by assigning an extra momenttun in random directions. with a rate and magnitude chosen to reproduce mass-loads and wind speeds similar to those observed.
In this model, gas particles are stochastically driven out of the dense star-forming medium by assigning an extra momentum in random directions, with a rate and magnitude chosen to reproduce mass-loads and wind speeds similar to those observed.
See Springel&IIernquist(2003a) for a detailed discussion of this method.
See \citet{SH02b} for a detailed discussion of this method.
Alost of our simulations employ a “strong” wind of μα»eed 484kms.+. but for the 10ἡ*\Ipe box (runs O3. P3. Q3. O4. Q5: collectively called ‘Q-series’) we also varied the wind strength.
Most of our simulations employ a “strong” wind of speed $484\,{\rm km\,s^{-1}}$, but for the $10\,h^{-1}{\rm Mpc}$ box (runs O3, P3, Q3, Q4, Q5; collectively called `Q-series') we also varied the wind strength.
Therefore. this Q-series can be used to study both the elfect of numerical resolution and the consequences of feedback. from galactic winds.
Therefore, this Q-series can be used to study both the effect of numerical resolution and the consequences of feedback from galactic winds.
The runs in the other simulation series then allow the extension of the strong wind results to smaller scales Clt-Series). or to larger box-sizes and hence lower redshift (CD- and "€i-Series).
The runs in the other simulation series then allow the extension of the strong wind results to smaller scales (`R-Series'), or to larger box-sizes and hence lower redshift (`D-' and `G-Series').
Our naming
Our naming
drops as one progresses outward through the fields: Stanek's Window (41%)). SWEEPS (39%)). Baade's Window )). and OGLE29 (35%)).
drops as one progresses outward through the fields: Stanek's Window ), SWEEPS ), Baade's Window ), and OGLE29 ).
Rich et ((2007) found no gradient in the bulge metallicity when comparing the spectroscopic metallicities of 17 M giants in the inner bulge to 14 M giants in Baade's window. but given their small sample. their result may be consistent with our findings.
Rich et (2007) found no gradient in the bulge metallicity when comparing the spectroscopic metallicities of 17 M giants in the inner bulge to 14 M giants in Baade's window, but given their small sample, their result may be consistent with our findings.
Zoccali et ((2008) found the bulge to decrease in metallicity along the minor axis beyond 4".
Zoccali et (2008) found the bulge to decrease in metallicity along the minor axis beyond $^{\rm o}$.
Note that our measured MDF variation is for the general population in each field. and cannot be investigated for the bulge in isolation. without proper-motion cleaning of the foreground stars.
Note that our measured MDF variation is for the general population in each field, and cannot be investigated for the bulge in isolation without proper-motion cleaning of the foreground stars.
We have used the TRILEGAL Galaxy model. with its default parameters (Girardi et 22005). to estimate that the foreground thin disk. thick disk. and halo together contribute of the total population in each field. with higher contamination at increasing. distance from the Galactic center.
We have used the TRILEGAL Galaxy model, with its default parameters (Girardi et 2005), to estimate that the foreground thin disk, thick disk, and halo together contribute of the total population in each field, with higher contamination at increasing distance from the Galactic center.
We performed à preliminary. proper-motion cleaning of the catalog in the SWEEPS field. by comparing our astrometry to that in Sahu et ((2006). and assuming the relative disk and bulge velocities of Clarkson et ((2008).
We performed a preliminary proper-motion cleaning of the catalog in the SWEEPS field, by comparing our astrometry to that in Sahu et (2006), and assuming the relative disk and bulge velocities of Clarkson et (2008).
This cleaning clearly reduces the presence of CMD features associated with the disk foreground (e.g.. the blue plume above the old MSTO). but the MDF in the remaining bulge population is not significantly changed from that shown in Figure 2.
This cleaning clearly reduces the presence of CMD features associated with the disk foreground (e.g., the blue plume above the old MSTO), but the MDF in the remaining bulge population is not significantly changed from that shown in Figure 2.
The combination of isochrones and synthetic spectra employed in Figure 2 was chosen because it spans nearly the full range of photometric indices in each field. demonstrating the nonlinear relationship between indices and [Fe/H].
The combination of isochrones and synthetic spectra employed in Figure 2 was chosen because it spans nearly the full range of photometric indices in each field, demonstrating the nonlinear relationship between indices and [Fe/H].
However. the zeropoint of the [Fe/H] scale is uncertain at the level of 0.2-0.3 dex. depending upon the actual extinction law in Baade's Window (changing Ry by 0.1 shifts the implied [Fe/H] by 0.1 dex). the assumed abundance pattern (e.g.. [a/Fe] as a function of [Fe/H]. particularly at high [Fe/H[). the photometric zeropoints. and the spectral library employed in the isochrone transformation.
However, the zeropoint of the [Fe/H] scale is uncertain at the level of 0.2–0.3 dex, depending upon the actual extinction law in Baade's Window (changing $R_V$ by 0.1 shifts the implied [Fe/H] by 0.1 dex), the assumed abundance pattern (e.g., $\alpha$ /Fe] as a function of [Fe/H], particularly at high [Fe/H]), the photometric zeropoints, and the spectral library employed in the isochrone transformation.
Putting the empirical ridge lines of our star clusters (322) in the same [zi] vs. [¢] plane would imply our assumed [Fe/H] scale underestimates the true [Fe/H] by ~0.2-0.3 dex down to [Fe/H] «—1.5. and that our metallicity scale 1s degenerate at lower metallicities.
Putting the empirical ridge lines of our star clusters 2) in the same $m$ ] vs. $t$ ] plane would imply our assumed [Fe/H] scale underestimates the true [Fe/H] by $\sim$ 0.2–0.3 dex down to [Fe/H] $\approx -1.5$, and that our metallicity scale is degenerate at lower metallicities.
The offset may be a real systematic error in our metallicity scale. but it may also be due to distinctions in the reddening law and/or abundance pattern between the cluster and bulge populations.
The offset may be a real systematic error in our metallicity scale, but it may also be due to distinctions in the reddening law and/or abundance pattern between the cluster and bulge populations.
The cluster photometry was not corrected for variations in extinction law in the same manner used for the bulge fields. because the MS locus for each cluster is very distinct from that in the bulge. given the single metallicity yet noisier photometry.
The cluster photometry was not corrected for variations in extinction law in the same manner used for the bulge fields, because the MS locus for each cluster is very distinct from that in the bulge, given the single metallicity yet noisier photometry.
Our SWEEPS field includes 13 of the l6 candidate exoplanet hosts found in theHST transit survey of Sahu et (2006).
Our SWEEPS field includes 13 of the 16 candidate exoplanet hosts found in the transit survey of Sahu et (2006).
Two of these 13 candidates are too faint and red to appear in our C images. but the remaining I] have photometry enabling their placement in the |v] vs. [t] plane for the general SWEEPS population (Figure 2: diamonds).
Two of these 13 candidates are too faint and red to appear in our $C$ images, but the remaining 11 have photometry enabling their placement in the $m$ ] vs. $t$ ] plane for the general SWEEPS population (Figure 2; ).
Two of these 11 candidates have radial velocities that support their planetary nature. and are highlighted in Figure 2 diamonds).
Two of these 11 candidates have radial velocities that support their planetary nature, and are highlighted in Figure 2 ).
Although the zeropoint for our assumed [Fe/H] scale Is uncertain at the level of —0.2-0.3 dex (see $33.3). the relative [Fe/H] measurements for the exoplanet hosts and the general population in the SWEEPS field are much more secure. because the same systematic uncertainties apply to both.
Although the zeropoint for our assumed [Fe/H] scale is uncertain at the level of $\sim$ 0.2–0.3 dex (see 3.3), the relative [Fe/H] measurements for the exoplanet hosts and the general population in the SWEEPS field are much more secure, because the same systematic uncertainties apply to both.
It is clear from Figure 2 that the candidate exoplanet hosts predominantly fall in the metal-rich end of the bulge MDF - a population that is already skewed toward high metallicity.
It is clear from Figure 2 that the candidate exoplanet hosts predominantly fall in the metal-rich end of the bulge MDF – a population that is already skewed toward high metallicity.
Aside from a single candidate at the metal-poor end of the distribution. the remaining 10 candidates are more metal-rich than half the population. with 7 1n the top quartile.
Aside from a single candidate at the metal-poor end of the distribution, the remaining 10 candidates are more metal-rich than half the population, with 7 in the top quartile.
A KS test of the implied metallicities in the exoplanet hosts and general population indicates that the chance they are both drawn from the same parent population ts less than2%.
A KS test of the implied metallicities in the exoplanet hosts and general population indicates that the chance they are both drawn from the same parent population is less than.
. We have performed a preliminary analysis of the data from our WFC3 Galactic Bulge Treasury Program.
We have performed a preliminary analysis of the data from our WFC3 Galactic Bulge Treasury Program.
Keeping in mind the various systematic uncertainties at this stage. our analysis of the dwarf stars supports the picture of the bulge gleaned from investigations of the brighter giant stars (see Zoccah 2010 and references therein).
Keeping in mind the various systematic uncertainties at this stage, our analysis of the dwarf stars supports the picture of the bulge gleaned from investigations of the brighter giant stars (see Zoccali 2010 and references therein).
Our IR photometry reaches the knee on the lower MS. and indicates that the population is predominantly old (~10 Gyr) in all of the bulge fields. with no obvious age gradient.
Our IR photometry reaches the knee on the lower MS, and indicates that the population is predominantly old $\sim$ 10 Gyr) in all of the bulge fields, with no obvious age gradient.
The declining metallicities at increasing radius are seemingly inconsistent with the secular processes that are traditionally associated with the formation of a peanut-shaped bulge.
The declining metallicities at increasing radius are seemingly inconsistent with the secular processes that are traditionally associated with the formation of a peanut-shaped bulge.
Our findings are consistent with a classical bulge formed via rapid dissipative collapse (either monolithic or via the merger of independent components). but also consistent with a recently emerging formation paradigm. motivated by observations of gas-rich spirals at z2 (Genzel et 22005: Fórrster Schreiber et 22009).
Our findings are consistent with a classical bulge formed via rapid dissipative collapse (either monolithic or via the merger of independent components), but also consistent with a recently emerging formation paradigm, motivated by observations of gas-rich spirals at $z \sim 2$ (Genzel et 2008; Förrster Schreiber et 2009).
In this new paradigm. instabilities in gas-rich disks can drive early bulge formation over rapid timescales (e.g.. Immeli et 22004: Elmegreen et 22009).
In this new paradigm, instabilities in gas-rich disks can drive early bulge formation over rapid timescales (e.g., Immeli et 2004; Elmegreen et 2009).
Of the hundreds of extrasolar planets discovered to date. most have been found in the solar neighborhood via radial-velocity measurements.
Of the hundreds of extrasolar planets discovered to date, most have been found in the solar neighborhood via radial-velocity measurements.
A notable exception is the discovery of 16 candidate exoplanet hosts in the SWEEPS transit survey of the Galactic bulge (Sahu et 22006).
A notable exception is the discovery of 16 candidate exoplanet hosts in the SWEEPS transit survey of the Galactic bulge (Sahu et 2006).
Our multiband photometry of 11 of these hosts demonstrates that they fall almost exclusively at the high end of the MDF in this high-density. metal-rich field.
Our multiband photometry of 11 of these hosts demonstrates that they fall almost exclusively at the high end of the MDF in this high-density, metal-rich field.
Exoplanets in the distinct environment of the solar neighborhood are also found preferentially at high metallicity (e.g.. Fischer Valenti 2005).
Exoplanets in the distinct environment of the solar neighborhood are also found preferentially at high metallicity (e.g., Fischer Valenti 2005).
Out of the ~500 exoplanets discovered to date whose orbital periods range from a fraction of à day to well over five years. >100 have orbital periods less than five days. implying significant migration since formation.
Out of the $\sim$ 500 exoplanets discovered to date whose orbital periods range from a fraction of a day to well over five years, $>100$ have orbital periods less than five days, implying significant migration since formation.
The correlation of such planets with stars of high metallicity probably indicates that planets are preferentially formed in high-metallicity environments. or alternatively that planets migrate more easily under such conditions.
The correlation of such planets with stars of high metallicity probably indicates that planets are preferentially formed in high-metallicity environments, or alternatively that planets migrate more easily under such conditions.
Support for Program 11664 was provided by NASA through a grant from STSel. which is operated by AURA. Inc.. under NASA contract NAS 5-26555.
Support for Program 11664 was provided by NASA through a grant from STScI, which is operated by AURA, Inc., under NASA contract NAS 5-26555.
MZ acknowledges Fondecyt Regular 1085278.
MZ acknowledges Fondecyt Regular 1085278.
AR acknowledges ASI for support via the grant "COFIS-Analisi Dati.”
AR acknowledges ASI for support via the grant “COFIS-Analisi Dati.”
We appreciate useful discussions with J. Kalirat and A. Dotter.
We appreciate useful discussions with J. Kalirai and A. Dotter.
in excess of LOMAL. and a deliciency may be accounted for by a single massive galaxy at the center.
in excess of $10^{12}M_\odot$ and a deficiency may be accounted for by a single massive galaxy at the center.
Siuce the mass profile obtained by Broadhurstetal.(2005b) includes the contribution of galaxies. we should at least examine the possible effect to our models dealiug solely with particles.
Since the mass profile obtained by \citet{TB05B} includes the contribution of galaxies, we should at least examine the possible effect to our models dealing solely with particles.
Qur tov galaxy has a total mass ofoy 2x10>A4. aud à Gaussian⋅ volume prolile⋅ with⋅ the scale lengtl oL 20 kpe.
Our toy galaxy has a total mass of $2\times10^{12}M_\odot$ and a Gaussian volume profile with the scale length of 20 kpc.
We did uot perform any fittiug or adjustment of the galaxy. model parameters. except that the total uass corresponds to the 3D eucircled mass deficieucy. with respect to the observec value.
We did not perform any fitting or adjustment of the galaxy model parameters, except that the total mass corresponds to the 3D encircled mass deficiency with respect to the observed value.
It turnecl out that this toy galaxy had πο effect on the particle mass. iv. although the EHSE was solved froi1 the origin as usual.
It turned out that this toy galaxy had no effect on the particle mass, $m$, although the EHSE was solved from the origin as usual.
Ou the otler hand. as can be seen iu the volume- and colui densities in Figs.ll and 16.. the model galaxy remedies the local mass celicteucy problem at the center.
On the other hand, as can be seen in the volume- and column densities in \ref{neu.rho} and \ref{neu.Sigma}, the model galaxy remedies the local mass deficiency problem at the center.
Iu Fie.18.. the Fermi levels Ep and classical kinetic energies Ey are plotted for the moclels without and with the tov egalaxy.
In \ref{neu.EfEk}, the Fermi levels $E_F$ and classical kinetic energies $E_K$ are plotted for the models without and with the toy galaxy.
The toy egalaxy acts as ai external fae)eravity source. whose effect on the 3D encircled mass is largest near the center.
The toy galaxy acts as an external gravity source, whose effect on the 3D encircled mass is largest near the center.
The gravity of the galaxy enliauces tle classical kiuetic energy of a particle aud clissolves the degeneracy.
The gravity of the galaxy enhances the classical kinetic energy of a particle and dissolves the degeneracy.
This is the reason why there is a dip of the Fermi level in the model with the galaxy.
This is the reason why there is a dip of the Fermi level in the model with the galaxy.
La other words. the EOS is softened locally by the stellar inass of the galaxy.
In other words, the EOS is softened locally by the stellar mass of the galaxy.
Ht should be emphasized that this is tre oulv if the galaxy 1s given as an exterbal lass xwce unallected by the dark matter distribution.
It should be emphasized that this is true only if the galaxy is given as an external mass source unaffected by the dark matter distribution.
Also this effect is pronounced ouly near the oriein of spherical svimuuetry. where the fractional contribution of the galaxy to the 3D encircled imass is tlie greatest.
Also this effect is pronounced only near the origin of spherical symmetry, where the fractional contribution of the galaxy to the 3D encircled mass is the greatest.
One may wonder how fermious at the cluster center could be degenerate.
One may wonder how fermions at the cluster center could be degenerate.
Here we first point out the fact that even relic neutrinos are moclerately degenerate. which used to be well kuowu (Weinberg1962).. but appears to be forgotten lately.
Here we first point out the fact that even relic neutrinos are moderately degenerate, which used to be well known \citep{Weinberg1962}, but appears to be forgotten lately.
There is a clear clistinetion between the black body of fermions aud that of yosous.
There is a clear distinction between the black body of fermions and that of bosons.
While the range of values of the Fermi distribution fuuctiou for one-particle sale is between O and 1. that of the Bose distribution function is between 0 aud x.
While the range of values of the Fermi distribution function for one-particle state is between 0 and 1, that of the Bose distribution function is between 0 and $\infty$.
So the fermio1 black body is partially degenerate. while the boson black body is not.
So the fermion black body is partially degenerate, while the boson black body is not.
Here we calculate the degree of degeneracy lor relic neurinos ininediately. after decoupling aud that after they become nonrelativistie. uuder the assumption of adiabatic expansion.
Here we calculate the degree of degeneracy for relic neutrinos immediately after decoupling and that after they become nonrelativistic, under the assumption of adiabatic expansion.
The current cosmology assumnes that the‘e are the equal utunbers of ueutrinos and lor each species.
The current cosmology assumes that there are the equal numbers of neutrinos and anti-neutrinos for each species.
So the present iuuuber deusity of each quantum statistically incepeucent species. n(0). is 112.6/2 = 56.3 7.
So the present number density of each quantum statistically independent species, $n(0)$, is 112.6/2 = 56.3 $^{-3}$.
At redshi[t z. the number cleusity n(z) is given by lumediately. after neutrino decoupling. when neutrinos were extremely relativistic. the Fermi level. By is given by
At redshift $z$, the number density $n(z)$ is given by Immediately after neutrino decoupling, when neutrinos were extremely relativistic, the Fermi level, $E_f$ is given by
Iu this section we shall analyze the global structure for the solutions fouud in the last section.
In this section we shall analyze the global structure for the solutions found in the last section.
Let us first consider the solutions for p—0.
Let us first consider the solutions for $p=0$.
From equation (251). we find that the corresponding metric can be written as
From equation \ref{Ip}) ), we find that the corresponding metric can be written as.
The eeoumetrical radius is giveu by τ-IS1).
The geometrical radius is given by =lS_0(-t).
Without loss of generality. we asstune that Sy>0.
Without loss of generality, we assume that $S_0>0$.
Since A. is à function off only. it is easy to sec that Ry is alwavs positive. since fF>0.
Since ${\cal R}$ is a function of $t$ only, it is easy to see that ${\cal R}_{,t}$ is always positive, since $-t > 0$.
The whole spacetime is trapped. as one can sce from the outeoiug and ingoing null ecodesics expansions. which arc now eiveu by the equations (see Appendix À)
The whole spacetime is trapped, as one can see from the outgoing and ingoing null geodesics expansions, which are now given by the equations (see Appendix A)
the seed photon source to the scattering electron population. i.c. the seed photons must first interact withabove the jet not at its base but thousands of gravitational radii the disc.
the seed photon source to the scattering electron population, i.e. the seed photons must first interact with the jet not at its base but thousands of gravitational radii above the disc.
This geometry raises significant problems for explaining the energeties ancl variability of the emission (since the jet at that height would subtencd only a small solid angle as seen from the disc). not to mention the shape and strength of the ironAve line. which due to geometric ancl beaming cllects would be much narrower ancl weaker than observed. if the disc were illuminated by the jet from such a ereat height (e.g. Alarkoll&Nowak 2004)).
This geometry raises significant problems for explaining the energetics and variability of the emission (since the jet at that height would subtend only a small solid angle as seen from the disc), not to mention the shape and strength of the iron$K\alpha$ line, which due to geometric and beaming effects would be much narrower and weaker than observed, if the disc were illuminated by the jet from such a great height (e.g. \citealt{Markoff04}) ).
A likely explanation of the large mecdium-soft. lags observed. on time-scales of seconds. is that the lags are associated with the generation. of ancl propagation of accretion fluctuations within the inner regions of the disc before they reach the corona (Lyubarskit1997:Ixotov.Chu-razov.&Gillanoy2001:ArévaloUttley 2006).
A likely explanation of the large medium-soft lags observed on time-scales of seconds is that the lags are associated with the generation of and propagation of accretion fluctuations within the inner regions of the disc before they reach the corona \citep{Lyubarskii97,Kotov01,Arevalo06}.
. In fact. the observed lag behaviour on these anc sub-second time-scales is exactly that expected. from our earlier. clise-variahbility interpretation of the covariance spectrum. (Wilkinson&Uttley 2009).
In fact, the observed lag behaviour on these and sub-second time-scales is exactly that expected from our earlier disc-variability interpretation of the covariance spectrum \citep{Wilkinson09}.
. In. the present work. we have used. causal information to greatly strengthen that interpretation. so that aceretion instabilities can be firmly. identified as the source of X-ray variability. at least on time-scales of seconds.
In the present work, we have used causal information to greatly strengthen that interpretation, so that accretion instabilities can be firmly identified as the source of X-ray variability, at least on time-scales of seconds.
Furthermore. the unstable aceretion Dow is now shown be the standard. disc. with all the attendant: physical implications (see discussion in Wilkinson&Uttley. 20093). and not a hot optically thin flow.
Furthermore, the unstable accretion flow is now shown to be the standard disc, with all the attendant physical implications (see discussion in \citealt{Wilkinson09}) ), and not a hot optically thin flow.
The cise propagation mocel can explain the mecdium- lags. however the hard-medium lags - produced. where he disc blackbody does not contribute. to. the X-ray spectrum - still require explanation.
The disc propagation model can explain the medium-soft lags, however the hard-medium lags - produced where the disc blackbody does not contribute to the X-ray spectrum - still require explanation.
These lags may still » related to Compton upscattering of seed photons as the accretion [luctuations reach the base of the Lotjet. c.g. in a ivbrid model of disc propagation and the mock lteig.Ixv-alis.&Ciannios (2003).
These lags may still be related to Compton upscattering of seed photons as the accretion fluctuations reach the base of the jet, e.g. in a hybrid model of disc propagation and the model of \citet{Reig03}.
. Llowever. given the common £One requency-dependence of the low-frequency lags which seers o be independent of the energy-bands chosen. it seems likely hat the lags at hard energies are directly related: to the same underlving propagation mechanism. which procduces he meclitmesolt lags.
However, given the common $\nu^{-0.7}$ frequency-dependence of the low-frequency lags which seems to be independent of the energy-bands chosen, it seems likely that the lags at hard energies are directly related to the same underlying propagation mechanism which produces the medium-soft lags.
For example. lags at harder energies can be simply produced if the mass-accretion Iuctuations in the disc. prelerentially generate softer and. then harder power-law emission as they propagate inwards.
For example, lags at harder energies can be simply produced if the mass-accretion fluctuations in the disc preferentially generate softer and then harder power-law emission as they propagate inwards.
Such effects could be produced if there is a soft coronal component above he cise while the hard emission is produced. centrally (e.g. Ixotov.Churazoyv.&CGilfanov200ArévaloUttley2
Such effects could be produced if there is a soft coronal component above the disc while the hard emission is produced centrally (e.g. \citealt{Kotov01,Arevalo06}) ).
006) The switch in lag behaviour on short time-scales suggests that the disc lags the power-law variations by a ew ms on these time-scales. consistent with the light-travel ags expected from a power-law Component separated from. he disc by only tens of gravitational radii at most.
The switch in lag behaviour on short time-scales suggests that the disc lags the power-law variations by a few ms on these time-scales, consistent with the light-travel lags expected from a power-law component separated from the disc by only tens of gravitational radii at most.
This ag signature represents the first evidence in DIINIUDs for a disc thermal reverberation lag.
This lag signature represents the first evidence in BHXRBs for a disc thermal reverberation lag.
This is the time-lae cue o the light travel-time from the central power-law emitting region to the disc where the hard. N-ravs are reprocessed into blackbody emission.
This is the time-lag due to the light travel-time from the central power-law emitting region to the disc where the hard X-rays are reprocessed into blackbody emission.
In order to use these lags to map he disc. e.g. to constrain the inner radius. we would. need o make assumptions about the geometry. of the power-aw emitting region and the emissivitv. profile of the clise Mackbods.
In order to use these lags to map the disc, e.g. to constrain the inner radius, we would need to make assumptions about the geometry of the power-law emitting region and the emissivity profile of the disc blackbody.
: “Phis modelling effort is beyond the scope of this work.
This modelling effort is beyond the scope of this work.
Finally. it is important to note that although the lags observed for variations on time-scales of less than a second appear to be caused by X-ray heating of the disc. this does not necessarily imply that the variability on these shorter time-scales is not also generated in the disc.
Finally, it is important to note that although the lags observed for variations on time-scales of less than a second appear to be caused by X-ray heating of the disc, this does not necessarily imply that the variability on these shorter time-scales is not also generated in the disc.
An accretion instability at small cise radii will modulate only disc emission inside that radius (e.g. Arevalo&Uttley2006)). so the intrinsic variability in disc emission. will be small and likely to be dominated by X-ray heating elfects when the mass lluctuations reach the central power-law emitting region.
An accretion instability at small disc radii will modulate only disc emission inside that radius (e.g. \citealt{Arevalo06}) ), so the intrinsic variability in disc emission will be small and likely to be dominated by X-ray heating effects when the mass fluctuations reach the central power-law emitting region.
We would like to thank the anonvmous referee for valuable comments.
We would like to thank the anonymous referee for valuable comments.
PU js supported. by an SPEC Advanced Fellowship and TW is supported by an SPEC postgraduate studentship grant.
PU is supported by an STFC Advanced Fellowship and TW is supported by an STFC postgraduate studentship grant.
The research. loading to these results has received. funding from the European Communitys Seventh Framework Programme (FDP7/2007-2013) uneler erant agreement number LEN 215212 “Black Hole Universe?
The research leading to these results has received funding from the European Community's Seventh Framework Programme (FP7/2007-2013) under grant agreement number ITN 215212 “Black Hole Universe”.
This work was partly funded. by the Buncdesministerium fine Wirtschaft ancl Technologie through Deutsches Zentrum fur Luft- und Raumlahrt grants 50 OR. 0701 and 50 Ol OSOS.
This work was partly funded by the Bundesministerium fürr Wirtschaft and Technologie through Deutsches Zentrum fürr Luft- und Raumfahrt grants 50 OR 0701 and 50 OR 0808.
"This work is based on observations obtained with AAMA-Newlon. an ESA science mission. with instruments and contributions directly. Funded by LSA Member States and NASA.
This work is based on observations obtained with , an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA.
to explore completely the parameter space of each model and to lind the largest possible class of solutions Chat is consistent with the data.
to explore completely the parameter space of each model and to find the largest possible class of solutions that is consistent with the data.
This determines (he structure of the parameter space in detail. which would illuminate any other less obvious model degeneracies (hat mieht exist.
This determines the structure of the parameter space in detail, which would illuminate any other less obvious model degeneracies that might exist.
It is useful to enumerate (he parameters of each of (he 3 models discussed in §2 before discussing the results of our parameter space search.
It is useful to enumerate the parameters of each of the 3 models discussed in $\S$ \ref{sec:theory} before discussing the results of our parameter space search.
There is a single dimensionless theoretical parameter C for the truncated. Ostriker model. and 3 dimensionless theoretical parameters for each of the magnetic models: ο.1.8 [or the GS models and. C.D..D, for ihe FP model.
There is a single dimensionless theoretical parameter $C$ for the truncated Ostriker model, and 3 dimensionless theoretical parameters for each of the magnetic models: ${C,\beta,\theta}$ for the GS models and ${C,\Gz,\Gphi}$ for the FP model.
All models require 5 additional parameters Chat reflect various observational complexities and unknown quantities.
All models require 5 additional parameters that reflect various observational complexities and unknown quantities.
Two dimensional parameters p. and σ determine (he density scale ancl core radius.
Two dimensional parameters $\rhoc$ and $\sigma$ determine the density scale and core radius.