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Two of the remaining 19 sources fall outside the UVIS FoV, while the others are very close to the saturation limit in our catalogue or fall into the gaps of the UVIS mosaic. | Two of the remaining 19 sources fall outside the UVIS FoV, while the others are very close to the saturation limit in our catalogue or fall into the gaps of the UVIS mosaic. |
All these 67 sources had Ha excess emission at the time of the SB04 observations, but only 35 of them have Ha excess emission at the >5c level when our observations were taken. | All these 67 sources had $\alpha$ excess emission at the time of the SB04 observations, but only 35 of them have $\alpha$ excess emission at the $> 5\,\sigma$ level when our observations were taken. |
This number decreases to 23 if we consider only stars with EWua>10 AA. | This number decreases to 23 if we consider only stars with $_{H\alpha}
> 10$ . |
. When lowering the threshold to 3c, the number of stars with Ha excess in common with SB04 grows to 41, but only 25 of them have EWy,, >10A. | When lowering the threshold to $3\,\sigma$, the number of stars with $\alpha$ excess in common with SB04 grows to 41, but only 25 of them have $_{H\alpha}>$ 10. |
. This means that most %) of the sources showing Ha excess in the study of SB04 (March 1999) do not show it at the epoch of our observations (Aug 2009). | This means that most $\sim
65\,\%$ ) of the sources showing $\alpha$ excess in the study of SB04 (March 1999) do not show it at the epoch of our observations (Aug 2009). |
This is in line with the current understanding of the accretion mechanism whereby the in-falling material from the circumstellar disk is subject to bursts, corresponding to peaks in the Ha emission (e.g.Fernandezetal.1995). | This is in line with the current understanding of the accretion mechanism whereby the in-falling material from the circumstellar disk is subject to bursts, corresponding to peaks in the $\alpha$ emission \citep[e.g.][]{fer95}. |
. SB04 looked for matches between bright X-ray sources and objects with Ha excess emission by using archive Chandra X-Ray Observatory images. | SB04 looked for matches between bright X-ray sources and objects with $\alpha$ excess emission by using archive Chandra X-Ray Observatory images. |
Only ~1/3 of the 96 Ha excess emission sources in the WFPC2 data set have a ray emission. | Only $\sim1/3$ of the 96 $\alpha$ excess emission sources in the WFPC2 data set have a X-ray emission. |
We found that of the 23 stars with Ha excess emission in common between our catalogue and that of SB04, a total of 8 have X-ray emission, according to SB04. | We found that of the 23 stars with $\alpha$ excess emission in common between our catalogue and that of SB04, a total of 8 have X-ray emission, according to SB04. |
While it is expected PMS stars actively undergoing mass accretion to be detected in X-rays, it is also true that the physical processes related to X-ray flaring and those pertaining to mass accretion are not the same (seee.g.Feigelson2005). | While it is expected PMS stars actively undergoing mass accretion to be detected in X-rays, it is also true that the physical processes related to X-ray flaring and those pertaining to mass accretion are not the same \citep[see e.g.][]{fei05}. |
. As shown by SB04, we should not expect a complete match between the sources that are bright in X rays and those showing Ha excess at any given time. | As shown by SB04, we should not expect a complete match between the sources that are bright in X rays and those showing $\alpha$ excess at any given time. |
And, obviously, even less so when the observations are not simultaneous, as in the case of the Chandra and HST data used by there authors, because of the considerable variability to which the accretion process is subject (see above). | And, obviously, even less so when the observations are not simultaneous, as in the case of the Chandra and HST data used by there authors, because of the considerable variability to which the accretion process is subject (see above). |
The presence of stars with Ha excess emission the ZAMS in the CMD of Figure 6 offers newoverlapping support to the hypothesis of a spread in the age of the PMS stars in 33603 (see HAO8 and references therein), since this is the region where PMS stars older than MMyr are expected to be located. | The presence of stars with $\alpha$ excess emission overlapping the ZAMS in the CMD of Figure \ref{fig_ha_uv} offers new support to the hypothesis of a spread in the age of the PMS stars in 3603 (see HA08 and references therein), since this is the region where PMS stars older than Myr are expected to be located. |
One obvious issue to address is whether these objects are cluster members, as we know that there is a potentially significant contribution from field stars in this region (see Section 3.1)). | One obvious issue to address is whether these objects are cluster members, as we know that there is a potentially significant contribution from field stars in this region (see Section \ref{sec_err}) ). |
In order to investigate cluster membership for stars showing Ha emission, we can look at their spatial distribution compared to that of field stars. | In order to investigate cluster membership for stars showing $\alpha$ emission, we can look at their spatial distribution compared to that of field stars. |
Based on the distribution of Ha excess objects in the CMD, we define two regions containing bona-fide PMS stars (i.e. those with Ha excess of different ages). | Based on the distribution of $\alpha$ excess objects in the CMD, we define two regions containing bona-fide PMS stars (i.e. those with $\alpha$ excess of different ages). |
The MMyr PMS isochrone is used as a guide to define a rough separation between the young PMS population (<10 MMyr; hereafter YPMS) and the old PMS population (>10 MMyr; hereafter OPMS). | The Myr PMS isochrone is used as a guide to define a rough separation between the young PMS population $<10$ Myr; hereafter YPMS) and the old PMS population $>10$ Myr; hereafter OPMS). |
We finally define as field population all stars lying in the OPMS box and not showing Ha excess. | We finally define as field population all stars lying in the OPMS box and not showing $\alpha$ excess. |
The corresponding selection areas are shown as boxes in Figure 6.. | The corresponding selection areas are shown as boxes in Figure \ref{fig_ha_uv}. |
It is important to underline here that, having defined as bona-fide PMS stars only the objects with excess Ha emission, we are in practice setting a lower limit to the actual number of stars in the PMS phase in 33603, since some of them can be in the PMS stage without showing any Ha excess because of the variability of their Ha flux (see Section 4)). | It is important to underline here that, having defined as bona-fide PMS stars only the objects with excess $\alpha$ emission, we are in practice setting a lower limit to the actual number of stars in the PMS phase in 3603, since some of them can be in the PMS stage without showing any $\alpha$ excess because of the variability of their $\alpha$ flux (see Section \ref{sec_ha}) ). |
This is certainly the case for the largest majority of stars that lie to the right of the MMyr isochrone in 33. | This is certainly the case for the largest majority of stars that lie to the right of the Myr isochrone in 3. |
On the other hand, since we are interested in the radial distribution of PMS stars, selecting only objects with Ha excess emission does not affect the significance of our statistical analysis. | On the other hand, since we are interested in the radial distribution of PMS stars, selecting only objects with $\alpha$ excess emission does not affect the significance of our statistical analysis. |
The cumulative radial distribution of the four groups defined above (PMS, YPMS, OPMS and Field) with respect to the cluster center is shown in the left panel of Figure 7,, whereas in the right panel the positions of YPMS stars (crosses) and of OPMS stars (filled circles) are shown on the F656N band image. | The cumulative radial distribution of the four groups defined above (PMS, YPMS, OPMS and Field) with respect to the cluster center is shown in the left panel of Figure \ref{fig_ks}, whereas in the right panel the positions of YPMS stars (crosses) and of OPMS stars (filled circles) are shown on the $F656N$ band image. |
Tocalculate the radial distribution of these objects we the RA(J2000) = 1115"7:26 and DEC(J2000) = —61?adopted15/37"9 as coordinates for the cluster center, following SB04. | Tocalculate the radial distribution of these objects we adopted the RA(J2000) = $11^{\rm h}\,
15^{\rm m}\, 7\fs26$ and DEC(J2000) = $-61\degr\, 15\arcmin\, 37\farcs9$ as coordinates for the cluster center, following SB04. |
We exclude from our analysis the innermost 5" radius where there is a high concentration of | We exclude from our analysis the innermost $5\arcsec$ radius where there is a high concentration of |
By sampling at various instants we found that this ‘quasi-periodic’ radial movement of the SAL is maintained over at least 250 minutes (the entire statistical integration time). | By sampling at various instants we found that this `quasi-periodic' radial movement of the SAL is maintained over at least 250 minutes (the entire statistical integration time). |
We nole in addition that at certain (mes the superadiabatic peak in Procyon reaches twice the height of the solar superadiabatie peak. | We note in addition that at certain times the superadiabatic peak in Procyon reaches twice the height of the solar superadiabatic peak. |
When the outer part of the SAL is in an oplically thin region. the photons can more readilv carry away excess internal energv from the temperature (Inetuations. than they can in optically thick regions. | When the outer part of the SAL is in an optically thin region, the photons can more readily carry away excess internal energy from the temperature fluctuations, than they can in optically thick regions. |
As it looses its excess Internal enerev compared (o its surroundings at a greater rate. the fIuctuation is said to be radiativelv damped. | As it looses its excess internal energy compared to its surroundings at a greater rate, the fluctuation is said to be radiatively damped. |
Because of this damping it is reasonable to infer that the panode intensity amplitudes (which depend on the temperature fIuctuation) will be smaller when ihe SAL is in optically thinner regions. | Because of this damping it is reasonable to infer that the $p$ -mode intensity amplitudes (which depend on the temperature fluctuation) will be smaller when the SAL is in optically thinner regions. |
The root mean square vertical velocity is plotted as a function of log P in both the Sun and Procvon in Fig.5.. | The root mean square vertical velocity is plotted as a function of log P in both the Sun and Procyon in \ref{vzrms}. |
We see (hat the turbulent. velocity has a much lareer amplitudes in Procvon than in the Sun. | We see that the turbulent velocity has a much larger amplitudes in Procyon than in the Sun. |
The Procvon simulation predicts vertical velocities ancl velocity fluctuations in (he optically thin atmosphere 2-3 times larger (han in the Sun. | The Procyon simulation predicts vertical velocities and velocity fluctuations in the optically thin atmosphere 2-3 times larger than in the Sun. |
This is consistent with observation. | This is consistent with observation. |
A recent spectroscopic study of Proevon by Allende Prieto et al. ( | A recent spectroscopic study of Procyon by Allende Prieto et al. ( |
2002) concludes Chat a comparison of the velocily spans (in line bisectors) for the Sun and Procvon shows that "the span of Procyon's lines exceeds (he solar values by more than a [actor of 27. | 2002) concludes that a comparison of the velocity spans (in line bisectors) for the Sun and Procyon shows that “the span of Procyon's lines exceeds the solar values by more than a factor of 2”. |
A grav scale map of vertical velocities near optical depth unity. as previously shown in Fie. l.. | A gray scale map of vertical velocities near optical depth unity, as previously shown in Fig. \ref{29Mm-gran}, |
displav a pattern of granulation in Procvon similar to that observed in the Sun. but more chaotic and on a larger scale. | display a pattern of granulation in Procyon similar to that observed in the Sun, but more chaotic and on a larger scale. |
The lighter regions denote upllows. while (he darker regions al granule boundaries denote downllows. | The lighter regions denote upflows, while the darker regions at granule boundaries denote downflows. |
There is a difference of scale: in Proevon. the granule sizes average about 10.000. kin. while solar granules average about. 1200 km in horizontal scale. | There is a difference of scale; in Procyon, the granule sizes average about 10,000 km, while solar granules average about 1200 km in horizontal scale. |
A useful quantity is the ratio 2 of turbulent pressure D, to mean gas pressure 2. | A useful quantity is the ratio $R$ of turbulent pressure $P_{turb}$ to mean gas pressure ${\overline P}$. |
The overbar denotes a combined horizontal and temporal average. | The overbar denotes a combined horizontal and temporal average. |
This ratio can be written: which is a non-dimensional quantity. | This ratio can be written: which is a non-dimensional quantity. |
The quantities p and 0? are the average density and root mean square vertical velocity. defined as | The quantities $\rho$ and $v_z''$ are the average density and root mean square vertical velocity, defined as - ^2 |
Lere Z(8) is the observed temperature in the direction of the unit vector fi. 3;,,(8) is the spherical Harmonic function and Gf) is the Gabor window. | Here $T({\mathbf{\hat n}})$ is the observed temperature in the direction of the unit vector ${\mathbf{\hat n}}$, $Y_{\ell m}({\mathbf{\hat n}})$ is the spherical Harmonic function and $G({\mathbf{\hat n}})$ is the Gabor window. |
We now [ind an expression for the expectation value of Cy. | We now find an expression for the expectation value of $\tilde C_\ell$. |
We will here use a Gabor window which is azimuthallv svnimetric about à point Tio on the sphere. so that the window is only a function of the angular distance from this point on the sphere cos@=f:fi. | We will here use a Gabor window which is azimuthally symmetric about a point ${\mathbf{\hat n}}_0$ on the sphere, so that the window is only a function of the angular distance from this point on the sphere $\cos\theta={\mathbf{\hat n}}\cdot{\mathbf{\hat n}}_0$. |
Then one can write the Legendre expansion of the window as. One can also write. Inserting these two expressions in equation (6)) one gets where relation (B3)) for Wigner 3j Symbols were used. | Then one can write the Legendre expansion of the window as, One can also write, Inserting these two expressions in equation \ref{eq:psalm}) ) one gets where relation \ref{eq:wigy}) ) for Wigner 3j Symbols were used. |
Using this expression. the relation αμ)=Ciónó," and the orthogonality of Wigner symbols (equation B1)). one can write (653 as. With €, we will always mean (C)) when we are referring to the full sky. power spectrum. | Using this expression, the relation $\VEV{a_{\ell
m}^*a_{\ell'm'}}=C_\ell\delta_{\ell\ell'}\delta_{mm'}$ and the orthogonality of Wigner symbols (equation \ref{eq:wigort}) ), one can write $\VEV{\tilde C_\ell}$ as, With $C_\ell$ we will always mean $\VEV{C_\ell}$ when we are referring to the full sky power spectrum. |
In this expression. Af.7) is the Gabor kernel. The Legendre coellicients g;. are found by the inverse Legendre transformation. where @ is the cut-olf angle where the window goes to zero. | In this expression, $K(\ell,\ell')$ is the Gabor kernel, The Legendre coefficients $g_\ell$, are found by the inverse Legendre transformation, where $\theta_\mathrm{C}$ is the cut-off angle where the window goes to zero. |
One sees from the expression for the kernel. that there is no dependency on fio. | One sees from the expression for the kernel, that there is no dependency on ${\mathbf{\hat n}}_0$. |
This means that £653 is the same. independent on where the Gabor window is centred. | This means that $\VEV{\tilde C_\ell}$ is the same, independent on where the Gabor window is centred. |
In the rest of this section we will study the shape of this kernel which couples the €, on the apocdised sphere. | In the rest of this section we will study the shape of this kernel which couples the $\tilde C_\ell$ on the apodised sphere. |
In Fig. | In Fig. |
1 we have plotted the kernel for a Gaussian Gabor window. with 5 and 15 degrees ENIEIM. (corresponding to @=2.12 and @= 6.387) and @=30. | \ref{fig:kernelg} we have plotted the kernel for a Gaussian Gabor window, with $5$ and $15$ degrees FWHM (corresponding to $\sigma=2.12^\circ$ and $\sigma=6.38^\circ$ ) and $\theta_\mathrm{C}=3\sigma$. |
One sees that the kernel is centred about (=(5. and falls off rapidly. | One sees that the kernel is centred about $\ell=\ell'$, and falls off rapidly. |
Eig. (2)) | Fig. \ref{fig:kernelth}) ) |
shows the same for the corresponding top-hat Gabor windows. The top-hat windows are covering the same area on the sky as the corresponding Gaussian windows in Fig. | shows the same for the corresponding top-hat Gabor windows, The top-hat windows are covering the same area on the sky as the corresponding Gaussian windows in Fig. |
1. (Gc is the same). | \ref{fig:kernelg}
$\theta_\mathrm{C}$ is the same). |
Ones sees that the diagonal is broader for the smaller windows indicating stronger couplings. | Ones sees that the diagonal is broader for the smaller windows indicating stronger couplings. |
Another thing to notice is that whereas the kernel for the top-hat Gabor window onlv falls bv about 4 orders of magnitude [rom the diagonal to the fav oll-diagonal elements. the Gaussian Gabor kernel falls by about 8 orders of magnitude (the vertical axis on the four plots are the same). | Another thing to notice is that whereas the kernel for the top-hat Gabor window only falls by about 4 orders of magnitude from the diagonal to the far off-diagonal elements, the Gaussian Gabor kernel falls by about 8 orders of magnitude (the vertical axis on the four plots are the same). |
The smooth cut-olf of the Gaussian Gabor window cuts olf long range correlations in spherical harmonic space. | The smooth cut-off of the Gaussian Gabor window cuts off long range correlations in spherical harmonic space. |
One of the aims of the first part of this paper is to see how the pseudo power spectrum of a given disc on the sky. (top-hat window) is alfected by the multiplication with a Gaussian Gabor window. | One of the aims of the first part of this paper is to see how the pseudo power spectrum of a given disc on the sky (top-hat window) is affected by the multiplication with a Gaussian Gabor window. |
For this reason the pseudo spectrum will be studied for à top-hat and a Gaussian covering the same area on the sky. | For this reason the pseudo spectrum will be studied for a top-hat and a Gaussian covering the same area on the sky. |
We will also study a top-hat window which has the same integrated area as the Gaussian window. | We will also study a top-hat window which has the same integrated area as the Gaussian window. |
The cut-off angle θε=6i for these windows is given by | The cut-off angle $\theta_\mathrm{C}=\theta_\mathrm{int}$ for these windows is given by |
Ilieh Velocity Clouds (IVCs) (Braun Burton 1999: Blitz et al. | High Velocity Clouds (HVCs) (Braun Burton 1999; Blitz et al. |
1998). | 1998). |
Tn our scenario. it would be possible to accomut or sone of these objects provided we modify the model of 822.2. | In our scenario, it would be possible to account for some of these objects provided we modify the model of 2.2. |
We currently asse that there is no accretion after reionization. and that most of the eas that accretes )ofore τνς 1s converted mto stars. | We currently assume that there is no accretion after reionization, and that most of the gas that accretes before $\zre$ is converted into stars. |
Towever. it may he that accretion starts again at low redshift. after the level of the UVbackground drops (e.g...Babul Rees 1992: IKepuor. Babul Spereel 1997). aud these late-accretiug svstenis might orl stars iuefficieutlv aud retain their gas as IIT. | However, it may be that accretion starts again at low redshift, after the level of the UV background drops (e.g., Babul Rees 1992; Kepner, Babul Spergel 1997), and these late-accreting systems might form stars inefficiently and retain their gas as HI. |
Tf this is he case. IIVCs av be associated with subhalos that are accreted at late times. | If this is the case, HVCs may be associated with subhalos that are accreted at late times. |
We fiud that. on average. ~GO% of surviving subhlalos fall iuto the host halo after 7=1. aud ~10 fall in after +=0.5. | We find that, on average, $\sim 60\%$ of surviving subhalos fall into the host halo after $z=1$, and $\sim 40 \%$ fall in after $z=0.5$. |
Another prediction is that there should be a diffuse stellar distribution in the Milkv. Way halo associated with the disvuption of many galactic satellites (sce Figure 1). | Another prediction is that there should be a diffuse stellar distribution in the Milky Way halo associated with the disruption of many galactic satellites (see Figure 1). |
Tf we asstune that the cestroved subhalos had the same stellar couteut as tho surviving halos. then we can estimate the radial density profile of this component by placing the stars from each disrupted halo at the apoceuter of its orbit. where they would spend most of them time. | If we assume that the destroyed subhalos had the same stellar content as the surviving halos, then we can estimate the radial density profile of this component by placing the stars from each disrupted halo at the apocenter of its orbit, where they would spend most of their time. |
This calculation vields a density profile p«(r)xk&P". with a=2540.3. extending from rz10αυfFσ]ηρο, which is roughly consistent with the distribution of known stellar halo populations such as RR Lyrac variables (6.8. Wetterer AIicCraw 1996). | This calculation yields a density profile $\rho_*(r) \propto
r^{-\alpha}$, with $\alpha = 2.5 \pm 0.3$, extending from $r \simeq 10
- 100 \hkpc$, which is roughly consistent with the distribution of known stellar halo populations such as RR Lyrae variables (e.g., Wetterer McGraw 1996). |
The normalization of the profile is more uucertain. but for the parameters n4=8 and f=0.3. hasand with the assumption that each halo with .... a nuass in stars of M.=f(Q,/OQje. 1. we fud that the stellar mass of the disrupted. component is AL,~5οTAL... | The normalization of the profile is more uncertain, but for the parameters $\zre = 8$ and $f = 0.3$, and with the assumption that each halo with $z_f > \zre$ has a mass in stars of $M_{*} = f (\Omega_b/\Omega_0) \epsilon_* M_{a}$ , we find that the stellar mass of the disrupted component is $M_*
\sim 5 \times 10^{8} \hMsun$. |
This diffuse distribution could make up a large fraction of the stellar halo. perhaps all of it. | This diffuse distribution could make up a large fraction of the stellar halo, perhaps all of it. |
Observationallv. it may be difficult to distinguish a dizupted populaion from a stellar alo formed by other means. but perhaps phase space substructure may provide a useful diagnostic (6.8.. Johustou 1998: Ποιά et al. | Observationally, it may be difficult to distinguish a disrupted population from a stellar halo formed by other means, but perhaps phase space substructure may provide a useful diagnostic (e.g., Johnston 1998; Helmi et al. |
1999). | 1999). |
This disuptec population woulk not be present in models with suppressed small scale power or wari. dark iatter. | This disrupted population would not be present in models with suppressed small scale power or warm dark matter. |
However. it would be expected in the selEiuteracting dark matter scenario. | However, it would be expected in the self-interacting dark matter scenario. |
In this case. the distribution would probably extend to a larger radius because dark matter interactions would cistupt the satellite halos further «ot. | In this case, the distribution would probably extend to a larger radius because dark matter interactions would disrupt the satellite halos further out. |
There are other problems facing the CDM. livypothesis. such as the posside disagreement between the predicted inner slopes o halo profiles aud tιο rotation curves of dwirf aud LSB ealaxi105 (Moore et al. | There are other problems facing the CDM hypothesis, such as the possible disagreement between the predicted inner slopes of halo profiles and the rotation curves of dwarf and LSB galaxies (Moore et al. |
1999: RIravtsov et al | 1999; Kravtsov et al. |
1998: Flores Pruuack 1991: Moore 1991). | 1998; Flores Primack 1994; Moore 1994). |
The uechanisui proposed here does not solve this problem. hough more complicated effects of eas cynics aud star ornmation nüeht do so. | The mechanism proposed here does not solve this problem, though more complicated effects of gas dynamics and star formation might do so. |
We have shown that one of the woblems facing CDM. cau be resolved by a simple gas (lenuandeal mechanisin. | We have shown that one of the problems facing CDM can be resolved by a simple gas dynamical mechanism. |
If this solution is the right one. hen the dark matter structure of the Milkv Way halo rescibles a scaled-down version of a vpical galaxy cluster. mt most of the low-11ass Milky. Way subhalos formed too ate to accrete eas and become observable dsvarf ealaxies. | If this solution is the right one, then the dark matter structure of the Milky Way halo resembles a scaled-down version of a typical galaxy cluster, but most of the low-mass Milky Way subhalos formed too late to accrete gas and become observable dwarf galaxies. |
We thank Andrew Could aud Jordi Miralda-Escudé or useful discussions. | We thank Andrew Gould and Jordi Miralda-Escudé for useful discussions. |
This work was supported in part w NASA LTSA eraut NACGS5S-3525 and NSF eraut. AST-ςJNU2568. | This work was supported in part by NASA LTSA grant NAG5-3525 and NSF grant AST-9802568. |
Support for A.WAN. was provided by NASA hrough Hubble Fellowship erant HE-01121.01-09À. from he Space Telescope Science Iustitute. which is operated by he Association of Universities for Research in Astrououy, Iuc.. uuder NASAcontract NÀS5-26555. | Support for A.V.K. was provided by NASA through Hubble Fellowship grant HF-01121.01-99A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASAcontract NAS5-26555. |
The stellar population mixtures in 3 of the 4 Itacdio Galaxies is similar to that of ellipticals. which is expected. since their host galaxies are ellipticals. | The stellar population mixtures in 3 of the 4 Radio Galaxies is similar to that of ellipticals, which is expected, since their host galaxies are ellipticals. |
Phere is some difference in the contribution from stars of 100. Myr. which is larger in Raclio Galaxies than in ellipticals. and in the contribution [rom stars with | Cyr and stars with LO Cir and Z/Z.=0. which is smaller than in ellipticals. | There is some difference in the contribution from stars of 100 Myr, which is larger in Radio Galaxies than in ellipticals, and in the contribution from stars with 1 Gyr and stars with 10 Gyr and $_{\odot}=0$, which is smaller than in ellipticals. |
PISSO745-19 shows cdillerent results. with a smaller contribution of 10 Gyr. high Z/Z. stars and larger contributions of the vounger populations. which could be due to the fact that it is in the middle of a cooling How. (Cardiel. Gorgas Aragonn-Salamanca 1995). | PKS0745-19 shows different results, with a smaller contribution of 10 Gyr, high $_{\odot}$ stars and larger contributions of the younger populations, which could be due to the fact that it is in the middle of a cooling flow (Cardiel, Gorgas Arag\'onn-Salamanca 1995). |
The results presented in the previous section are in sharp contrast with the traditional view WKhoski 1978) that the spectra of Sevfert 2s are essentially composed. of an old stellar population plus an underlying FC. | The results presented in the previous section are in sharp contrast with the traditional view Koski 1978) that the spectra of Seyfert 2s are essentially composed of an old stellar population plus an underlying FC. |
We have shown tmt: (1) Sevfert 28 present a wide breadth of stellar population characteristics ancl thus cannot be adequately represented by a single starlight template. ( | We have shown that: (1) Seyfert 2s present a wide breadth of stellar population characteristics and thus cannot be adequately represented by a single starlight template. ( |
2) There are substantial «iferences. between the stellar populations of σον[ος 2s ancl elliptical galaxies. particularly regarding the contribution from stars with LO Gyr and Z/Z.=0.6 and stars with LOO Myr. ( | 2) There are substantial differences between the stellar populations of Seyfert 2s and elliptical galaxies, particularly regarding the contribution from stars with 10 Gyr and $_{\odot}=0.6$ and stars with 100 Myr. ( |
3) The contribution of age <10 AlAIvr stars and an FC is small. rarely exceeding 10 per cent of the light at AbSTOA.. | 3) The contribution of age $\le 10$ Myr stars and an FC is small, rarely exceeding 10 per cent of the light at $\lambda$. |
Of particular interest in our analvsis are the results for Mrk1210. ΑΙκος. Mrk573.. Mrk607.. NCC'1358. and ος20, | Of particular interest in our analysis are the results for Mrk1210, Mrk348, Mrk573, Mrk607, NGC1358 and 3C33. |
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