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We constrain the ft so that the two absorption features on oue hand. aud the two enüssion features on the other. have the same EWIIM. and are separated by 9 (as is expected from the doublet cousponeuts AAL393.1102). | We constrain the fit so that the two absorption features on one hand, and the two emission features on the other, have the same FWHM and are separated by 9 (as is expected from the doublet components $\lambda\lambda$ 1393,1402). |
The results of the fit are presented in Table 2. | The results of the fit are presented in Table 2. |
The FWIIM of the cmission line features are consistent with the value measured for the CTV aud Lyo emission. | The FWHM of the emission line features are consistent with the value measured for the CIV and $\alpha$ emission. |
The absorption feature presents a much larger FWIIM than | The absorption feature presents a much larger FWHM than |
over à powerlaw at low energies. the derived: parameters are very similar throughout Models D. E EF. Similarly for SWIFT J2127.4315654. the parameters obtained with these three models are all consistent and indeed the resulting spin parameter is also consistent with previous findings by Aliniutti et al. ( | over a powerlaw at low energies, the derived parameters are very similar throughout Models D, E F. Similarly for SWIFT J2127.4+5654, the parameters obtained with these three models are all consistent and indeed the resulting spin parameter is also consistent with previous findings by Miniutti et al. ( |
2009) although they find a much higher emissivity index is required (q=5.3F1 compared to g=2.2"n from Model E). | 2009) although they find a much higher emissivity index is required $q=5.3^{+1.7}_{-1.4}$ compared to $q=2.2^{+0.3}_{-0.9}$ from Model E). |
As discussed. previously. the spin constraint obtained here for Fairall 9 agrees with that found by Sehmoll et al. ( | As discussed previously, the spin constraint obtained here for Fairall 9 agrees with that found by Schmoll et al. ( |
2009). but onlv in the case where they ignore the spectrum. below 2kkeV (α—0.5Py compared to a—0.44.ft found here in Model D). | 2009), but only in the case where they ignore the spectrum below keV $a=0.5^{+0.1}_{-0.3}$ compared to $a=0.44^{+0.04}_{-0.11}$ found here in Model D). |
According to Model D. an intermediate spin of the central black hole within NGC 7469 is found. a—0.69.AUIDros. | According to Model D, an intermediate spin of the central black hole within NGC 7469 is found, $a=0.69^{+0.09}_{-0.09}$. |
This cpis also consistent. with. that found⋅ withinA . ↳∖↓⋯⇂∢⊾↓↓⊲↙∣∶∪⋅⋀−≽↕⋮⊽⊥∣↕⊳⊳∖⊔⋏∙≟⋏∙≟∢⊾⊳∖⊔⊔⋏∙≟↿⇂↥⋜∐↿↓⊔⋅⊳∖↓≻↓⊔∪⇂↿↓⊔⊳∖ayOL : . ⋅⋠ object is indeed e0.7. | This is also consistent with that found within Model F $a=0.72^{+0.18}_{-0.17}$, suggesting that the spin of this object is indeed $a\sim0.7$. |
‘These findings therefore suggest that the two cillering interpretations give concurring spin constraints only when the blurred reflector is not required to model the soft excess. leading to the conclusion that the component responsible for the soft. excess must. be. independently fitted before conclusive spin constraints can be made. | These findings therefore suggest that the two differing interpretations give concurring spin constraints only when the blurred reflector is not required to model the soft excess, leading to the conclusion that the component responsible for the soft excess must be independently fitted before conclusive spin constraints can be made. |
Even so. the origin of the soft excess is not currently known. e.g. see Cierlinski Done (2004). | Even so, the origin of the soft excess is not currently known, e.g. see Gierlinski Done (2004). |
The relatively constant temperature. of the soft excess versus DII mass suggests that it may not arise [rom cürect thermal emission from the accretion disc. since the accretion temperature properties should scale with Moy1/4 ina standard. accretion disc. | The relatively constant temperature of the soft excess versus BH mass suggests that it may not arise from direct thermal emission from the accretion disc, since the accretion temperature properties should scale with $M_{\rm BH}^{-1/4}$ in a standard accretion disc. |
An atomic origin of the soft excess has been suggested. however no obvious spectral features are seen in high resolution data meaning that if atomic emission is responsible it must be significantly relativistically blurred. such as here in Model 5 (Itoss. Fabian Ballantyne 2002: Fiore. Matt Nicastro 1997). | An atomic origin of the soft excess has been suggested, however no obvious spectral features are seen in high resolution data meaning that if atomic emission is responsible it must be significantly relativistically blurred such as here in Model E (Ross, Fabian Ballantyne 2002; Fiore, Matt Nicastro 1997). |
Alternatively. it is suggested that the soft excess could arise from relativistically blurred absorption (Cierlinski Done 2004) from a cdilferentially rotating and outllowing disc wind (Murray Chiang 1997). | Alternatively, it is suggested that the soft excess could arise from relativistically blurred absorption (Gierlinski Done 2004) from a differentially rotating and outflowing disc wind (Murray Chiang 1997). |
The smeared absorption creates a smooth hole in the spectrum resulting in the apparent. soft excess and a hardening at high energies. | The smeared absorption creates a smooth hole in the spectrum resulting in the apparent soft excess and a hardening at high energies. |
Whilst in this paper we do not aim to determine the origin of the soft. excess. we find that a simple parametrization of the soft. excess continuum through a model such as provides a better fit to the spectra than an atomic origin from a highly blurred rellection component. | Whilst in this paper we do not aim to determine the origin of the soft excess, we find that a simple parametrization of the soft excess continuum through a model such as provides a better fit to the spectra than an atomic origin from a highly blurred reflection component. |
One possible alternate origin of the observed broad component in AGN spectra i8. through reflection of primary continuum photons olf a Compton thick disc wind. producing emission. absorption and broad. features in the Γον band (Sim et al. | One possible alternate origin of the observed broad component in AGN spectra is through reflection of primary continuum photons off a Compton thick disc wind, producing emission, absorption and broad features in the K band (Sim et al. |
2010). | 2010). |
Fast. outllows can add significant. complexity to the ehh region ancl can reproduce a wide range of spectral signatures owing to the many dilfering physical conditions. | Fast outflows can add significant complexity to the K region and can reproduce a wide range of spectral signatures owing to the many differing physical conditions. |
Sim ct al. ( | Sim et al. ( |
2010) note that such outllows can. however. have a significant awlect upon the soft X-ray spectrum: resulting in features such as highly blueshifted absorption lines. although these absorption are weaker than the features seen in the Why band. | 2010) note that such outflows can, however, have a significant affect upon the soft X-ray spectrum resulting in features such as highly blueshifted absorption lines, although these absorption are weaker than the features seen in the K band. |
The lines from a cise wind may not be observed. in objects such as those in this sample if they are observed relatively face-on. however even if we are not looking directly through the line of sight to the wind. it may be possible to observe broaclened emission. features in rellection. | The lines from a disc wind may not be observed in objects such as those in this sample if they are observed relatively face-on, however even if we are not looking directly through the line of sight to the wind, it may be possible to observe broadened emission features in reflection. |
For instance Sim et al. ( | For instance Sim et al. ( |
2010) predict broadened WI. emission down to ~5 kkeV in their disc wind spectra. with equivalent widths typically several tens of eV. Resulting from the work above on this sample of six bare Sevfert AGN. we conclude the following: | 2010) predict broadened K emission down to $\sim5$ keV in their disc wind spectra, with equivalent widths typically several tens of eV. Resulting from the work above on this sample of six 'bare' Seyfert AGN, we conclude the following: |
log [O -- log |O z 1.1, | log [O - log [O $\lesssim$ 1.4. |
Iu Fig. 8.. | In Fig. \ref{o3re}, |
top paucl. we compare the |O TH] line luninosity with the total radio Wuninosity at 178 ΑΠΣ. | top panel, we compare the [O III] line luminosity with the total radio luminosity at 178 MHz. |
WTEC are located in general at higher line hnuninositv with respect to LEC of simulay racio Γιάrosity. | HEG are located in general at higher line luminosity with respect to LEG of similar radio luminosity. |
This reflects their higher value of the Lroug/vLizx tatio already cliscussed | This reflects their higher value of the $L_{\rm [OIII]}/\nu L_{178}$ ratio already discussed in \ref{hegleg}. |
There are two notable excepions to this trend. namely 3C Os laud 3€ 3Tl. | There are two notable exceptions to this trend, namely 3C 084 and 3C 371. |
They are spectroscopically classified as LEG. but show au excess o fa factor of ~ 50 in line 5wission with respect to the sotrees of this class of sinilar radio power. | They are spectroscopically classified as LEG, but show an excess of a factor of $\sim$ 50 in line emission with respect to the sources of this class of similar radio power. |
Intriguingly. fjese are the wo sources with 1e highest core dominance ¢Xf the sample. with ratios P. τς of 0.69 and 0.33 respectivelv. | Intriguingly, these are the two sources with the highest core dominance of the sample, with ratios $P_{\rm core}$ / $L_{178}$ of 0.69 and 0.33 respectively. |
This implies that. by considering only their gentine extende radio enissjon. rev would be stronger outliers. reaching ratios of liue to radio cimission even hieher than observe lin IIEC. | This implies that, by considering only their genuine extended radio emission, they would be stronger outliers, reaching ratios of line to radio emission even higher than observed in HEG. |
While DEG are only fouid at radio hinunositics larger ian log Lass |erg/s] 232.8, LEG cover t1ο whole range of radio power covered by the suounuple of 3CR sources witli z-0.3. almost five orders of uaenitude roni log 30.7 to log Lyry~35.1. | While HEG are only found at radio luminosities larger than log $L_{178}$ [erg/s] $\gtrsim 32.8$, LEG cover the whole range of radio power covered by the subsample of 3CR sources with $<$ 0.3, almost five orders of magnitude from log $L_{178}\sim30.7$ to log $L_{178}\sim35.4$. |
There is a clear trend for increasing line huuimositv with radio power. as already foimd aid discussed by several authors aud as reported in tre Introduction. | There is a clear trend for increasing line luminosity with radio power, as already found and discussed by several authors and as reported in the introduction. |
However. we ure now in the position of considering separately the sub-populations of ρα aud LEC. | However, we are now in the position of considering separately the sub-populations of HEG and LEG. |
More quantitatively. we find that TEC οVv a ΠΜ correlation in the for: log LioILI] — 1.15 log Liz | 2.06. | More quantitatively, we find that HEG obey a linear correlation in the form: log $L_{\rm[O~III]}$ = 1.15 $~$ log $L_{178}$ + 2.96. |
The error in the slope is 0.11. wile the xius around the correlation is 0.13 dex. | The error in the slope is 0.11, while the rms around the correlation is 0.43 dex. |
We tested t1¢ possible iuflueuce of redshift in driving this correlation (oth quantities depend on 27) estimating the partial raik coefücieut. finding MOULino—0.L1. | We tested the possible influence of redshift in driving this correlation (both quantities depend on $z^2$ ) estimating the partial rank coefficient, finding $r_{{\rm [O III]} - L_{178},z}=0.41$. |
For 16 data-»omuts. the probability that this results from raucdomly distributed data is 0.005. | For 46 data-points, the probability that this results from randomly distributed data is $P = 0.005$ . |
For LEC we πα: log LoIH] = 0.99 log Liss | 7.65, | For LEG we find: log $L_{\rm[O~III]}$ = 0.99 $~$ log $L_{178}$ + 7.65. |
The error in fie slope is 0.09 and the rms around the correlation is 150 dex. | The error in the slope is 0.09 and the rms around the correlation is 0.50 dex. |
The partial rank coefficient IS HOULugo!151. corresponding to a probability of a random distribution of P?=0.002. | The partial rank coefficient is $r_{{\rm [O III]} - L_{178},z}=0.51$, corresponding to a probability of a random distribution of $P = 0.002$. |
The rela10115 (erived for WTEC and LEC differ by a factor of ~ | Üint15 conmnuon range of radio power. | The relations derived for HEG and LEG differ by a factor of $\sim$ 10 in the common range of radio power. |
The slopes of the liuc-radio huuinositv correlations of TEC aud LEG are insead consistent witlii the errors. | The slopes of the line-radio luminosity correlations of HEG and LEG are instead consistent within the errors. |
Considering iusead the Tine. we fineD rather simular resuts. as shown iu Fig. &.. | Considering instead the line, we find rather similar results, as shown in Fig. \ref{o3re}, |
1uiddle pane. | middle panel. |
The onu of the correlations are: log Ly, = 1.06 log Lyss | SLL (for TEC) aud log Ly, = 0.53 log Lary | 13.12 (for LEG). | The form of the correlations are: log $L_{\rm {H\alpha}}$ = 1.06 $~$ log $L_{178}$ + $\,\,\,$ 5.44 (for HEG) and log $L_{\rm {H\alpha}}$ = 0.83 $~$ log $L_{178}$ + 13.13 (for LEG). |
The main difference is the smaller offset between the two populations. reduced to a factor of ~ 3. | The main difference is the smaller offset between the two populations, reduced to a factor of $\sim$ 3. |
Also the | Also the |
For Model 1. ny can be computed as follows. | For Model 1, $n_{\rm s}$ can be computed as follows. |
The total number of photons emitted in a certain frequency interval over the duration of the flare is Αν)= NuPo)/ hv. | The total number of photons emitted in a certain frequency interval over the duration of the flare is $N_{\rm s}(\nu)\,=$ $\Delta t_{\rm obs}\,P_{\rm s}(\nu)\,/$ ${h\,\nu}$. |
The emitting particles occupy. a volume of 7R*. | The emitting particles occupy a volume of $\pi\,R^3$. |
Taking into account that the photon densitv rises [rom 0 to its final value as the emitting particles move through (he emission region. the lime averaged photon densitv is ↕∐⊔∐↲≺∢≀↧↪∖⊽≼↲∪↓≯⇀∖↕⋯⇂≼↲↥∃⋅≀↧↪∖⊽∐∐∐≀↧↴↕⋅≀↧↴↕⋅↖≺↽↔↴∏∐∐↲∐↥≸↽↔↴↕∖⇁≼↲⊳∖⇁≀↧↴∐≼↲↕⋅⊳∖⇁∪∐∐↲≀↧↴↕⋅↕⊔∐⊔≼↲∐≺∢∶ The factor (1—In2) arises [rom averaging the svnchrotron photon density over (he time (he emiltine particles (travel from d4 to ds. taking into account that the shell heieht increases from {1/2 to HE and the number of svnchrotron photons increases linearly from 0 (o its final value. | Taking into account that the photon density rises from 0 to its final value as the emitting particles move through the emission region, the time averaged photon density is In the case of Model 2, a similar argument gives after some arithmetic: The factor $(1-\ln{2})$ arises from averaging the synchrotron photon density over the time the emitting particles travel from $d_1$ to $d_2$, taking into account that the shell height increases from $H/2$ to $H$ and the number of synchrotron photons increases linearly from 0 to its final value. |
Finally. the time averaged fIuxes received al Earth can be computed frou: The optical depth per path length for 55.— € processes is computed with the equations of Gould&Schréder(1967) The first integral5 runs over the piteh angle5 distribution and the second integrates5 over the target photon frequencies. | Finally, the time averaged fluxes received at Earth can be computed from: The optical depth per path length for $\gamma\gamma\,\rightarrow$ $e^+\,e^-$ processes is computed with the equations of \citet{Gould1967}
The first integral runs over the pitch angle distribution and the second integrates over the target photon frequencies. |
The threshold frequency for pair creation reads and the pair-creation cross section is | The threshold frequency for pair creation reads and the pair-creation cross section is |
has detected no lens svstems with 9>4" (Phillips οἱ 22000). | has detected no lens systems with $\theta >4''$ (Phillips et 2000). |
This excludes the highly concentrated halosin w<—0.7 models as indicated by the 1o.90%.. and confidence contours in Fie. | This excludes the highly concentrated halosin $\omega\la -0.7$ models as indicated by the $1\sigma$, and confidence contours in Fig. |
2a. | 2a. |
We also plot the lo confidence levels that would be imposed if one wide separation lens svsteni were discovered. | We also plot the $1\sigma$ confidence levels that would be imposed if one wide separation lens system were discovered. |
We interpret the result. as indicating that higher large separation lensing rates would lead to more refined constraints in the w—ολ} plane. | We interpret the result as indicating that higher large separation lensing rates would lead to more refined constraints in the $\omega-c(M)$ plane. |
We can also compare the total number of lenses predicted by the models with the 18 systems found among ~12000 sources in JVAS/CLASS. | We can also compare the total number of lenses predicted by the models with the 18 systems found among $\sim 12000$ sources in JVAS/CLASS. |
We find the expected number of lenses with 8>0.3" (the approximate angular resolution of JVAS/CLASS) to be 20.7. 13.2. 15.5 and 9.6 for the ACDM. w=—2/3. —1/2. and —1/3 models shown in Fie. | We find the expected number of lenses with $\theta > 0.3''$ (the approximate angular resolution of JVAS/CLASS) to be $20.7$, $18.2$, $15.5$ and $9.6$ for the $\LCDM$, $\omega=-2/3$, $-1/2$, and $-1/3$ models shown in Fig. |
1. respectively. | 1, respectively. |
Because the SIS lensing cross section is several orders of magnitude higher than that of the NEW. the expected number of lenses is not sensitive to the concentration e(GM). | Because the SIS lensing cross section is several orders of magnitude higher than that of the NFW, the expected number of lenses is not sensitive to the concentration $c(M)$. |
Instead. it depends more stronely on the cooling mass M, because a larger M, allows more halos to be modeled as SIS. | Instead, it depends more strongly on the cooling mass $M_c$ because a larger $M_c$ allows more halos to be modeled as SIS. |
In Fig. | In Fig. |
2b we quantilv the dependence of the total lensing rate on M, and w by plotting the relative likelihood curves Lie)xp(c)(1—p())" for detecting AN,=18 lenses from N=12000 sources. where p(w) is the model-predicted lensing rate Lor a given w. | 2b we quantify the dependence of the total lensing rate on $M_c$ and $\omega$ by plotting the relative likelihood curves $L(\omega)\propto p(\omega)^{N_l} (1-p(\omega))^{N-N_l}$ for detecting $N_l=18$ lenses from $N=12000$ sources, where $p(\omega)$ is the model-predicted lensing rate for a given $\omega$. |
The range of M, is similar to that discussed in IXochanek White (2001). with ihe upper and lower limits close to the cooling masses used bv Porciani Macau (2000) ancl Li Ostriker (2001). | The range of $M_c$ is similar to that discussed in Kochanek White (2001), with the upper and lower limits close to the cooling masses used by Porciani Madau (2000) and Li Ostriker (2001). |
The 26 confidence level in Fig. | The $2\sigma$ confidence level in Fig. |
2b suggests that the constraint on the equation of state is not sensitive to moderate alterations to the cooling mass. | 2b suggests that the constraint on the equation of state is not sensitive to moderate alterations to the cooling mass. |
Models with iz—0.4 are disfavored in all three cases. | Models with $\omega \ga -0.4$ are disfavored in all three cases. |
Note (hat our calculations implicitly assume that each halo below AL. harbors sullicient barvons (o give an 915 prolile for thefofal mass. | Note that our calculations implicitly assume that each halo below $M_c$ harbors sufficient baryons to give an SIS profile for the mass. |
His likely. however. (hat lower mass halos be devoid of barvons and therelore retain a shallower profile (e.g.. NEW). | It is likely, however, that lower mass halos be devoid of baryons and therefore retain a shallower profile (e.g., NFW). |
Because an SIS lens has a much larger cross section (han an NEW lens of the same mass. modeling all low-mass halos as SIS would overestimate the total lensing rate. | Because an SIS lens has a much larger cross section than an NFW lens of the same mass, modeling all low-mass halos as SIS would overestimate the total lensing rate. |
Constraints on w derived under our assumptions are therefore conservative — if low-mass halos actually contribute smaller cross sections than our caleulations predict. more negative w would be needed to reproduce the observed lensing rate. | Constraints on $\omega$ derived under our assumptions are therefore conservative – if low-mass halos actually contribute smaller cross sections than our calculations predict, more negative $\omega$ would be needed to reproduce the observed lensing rate. |
Both large separation lensing probabilities ancl the total predicted number of lenses are allected by the cosmological parameters P and Q,,. | Both large separation lensing probabilities and the total predicted number of lenses are affected by the cosmological parameters $h$ and $\Omega_m$. |
Our investigations showthat varvine these parameters within the lo error limits given in Netterfield οἱ al. ( | Our investigations showthat varying these parameters within the $1\sigma$ error limits given in Netterfield et al. ( |
2001) leads to almost parallel shifts in the contour lines and likelihood curves in Fig. | 2001) leads to almost parallel shifts in the contour lines and likelihood curves in Fig. |
2 bv 0.075Hr or less in a. therefore not affecting the generality of our results. | 2 by $0.075$ or less in $\omega$ , therefore not affecting the generality of our results. |
As the figures show. the predicted lensing probability decreases as 10 changes [rom —1 towards 0. | As the figures show, the predicted lensing probability decreases as $w$ changes from $-1$ towards 0. |
The effect is particularly strong for lenses with @2 4". where the probability | The effect is particularly strong for lenses with $\theta\ga 4''$ , where the probability |
proportional to the mass of the satellite ancl mass density of the halo. | proportional to the mass of the satellite and mass density of the halo. |
However it depends. critically on the orbital structure of the halo because à particles orbit will respond most strongly near resonances between the satellites and ido particle's orbital frequencies. | However it depends critically on the orbital structure of the halo because a particle's orbit will respond most strongly near resonances between the satellite's and halo particle's orbital frequencies. |
Co-orbiting trajectories will have the strongest response but the total mass involved is small for the standard model. | Co-orbiting trajectories will have the strongest response but the total mass involved is small for the standard model. |
LHigher-order resonances will oe weaker but occur at smaller. ealactocentric radii where he mass density is high. | Higher-order resonances will be weaker but occur at smaller galactocentric radii where the mass density is high. |
The wake is the product of these competing ellects and. generally. the wake peaks far inside he satellite orbit. | The wake is the product of these competing effects and, generally, the wake peaks far inside the satellite orbit. |
As an example. Figures 2 and 3. show the space density distortions induced by the LMC orbit in the standard Kine model halo in the orbital plane. | As an example, Figures \ref{fig:halowake1} and \ref{fig:halowake2}
show the space density distortions induced by the LMC orbit in the standard King model halo in the orbital plane. |
The satellite has pericenter at P=Y and apocenter at 2=14 and here. orbits in the counter-clockwise direction. | The satellite has pericenter at $R=7$ and apocenter at $R=14$ and here, orbits in the counter-clockwise direction. |
Pericenter is VY=7.¥o0. | Pericenter is $X=7, Y=0$. |
If he satellite were completely outside the halo. the dipole response would be a linear displacement. representing the new center of mass position. | If the satellite were completely outside the halo, the dipole response would be a linear displacement representing the new center of mass position. |
Lhe wake would be proportional o. dp(r)/dyrο where e is the unit vector from the halo o the satellite center (Weinberg 1989)). | The wake would be proportional to $-d\rho(r)/d{\bf r}\cdot{\bf e}$ where ${\bf e}$ is the unit vector from the halo to the satellite center (Weinberg \nocite{Wein:89}) ). |
In. Figure 2.. the satellite is inside the halo. and the wake deviates from he pure displacement. | In Figure \ref{fig:halowake1}, the satellite is inside the halo, and the wake deviates from the pure displacement. |
The amplitude of the quadrupole (Fig. 3)) | The amplitude of the quadrupole (Fig. \ref{fig:halowake2}) ) |
is a [actor of roughly two smaller than the dipole. | is a factor of roughly two smaller than the dipole. |
The dominant. wake is near the satellite pericen as expected. but note the inner lobe of the wake at roughly idf the pericenter distance. | The dominant wake is near the satellite pericenter as expected, but note the inner lobe of the wake at roughly half the pericenter distance. |
This is due primarily to the resonance between satellite and halo orbital azimuthal requencies. | This is due primarily to the 2:1 resonance between satellite and halo orbital azimuthal frequencies. |
Although the relative density of this inner lobe is smaller than the primary outer one at. pericenter (phase ). both the proximity ancl spatial structure causes the force rom the inner wake to dominate over direct. force. from he satellite. | Although the relative density of this inner lobe is smaller than the primary outer one at pericenter (phase 0), both the proximity and spatial structure causes the force from the inner wake to dominate over direct force from the satellite. |
As the satellite approaches pericenter. (phase 32/2). these inner lobes become relatively stronger and can dominate the response. | As the satellite approaches pericenter (phase $3\pi/2$ ), these inner lobes become relatively stronger and can dominate the response. |
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