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The subject of the RCD-pt has been already widely debated in the literature (sec. e.g. Reuzini Duzzoui 1986. Corsi et al.
The subject of the RGB-pt has been already widely debated in the literature (see, e.g., Renzini Buzzoni 1986, Corsi et al.
1991. €drardi Bertelli 1998 au references therein) aud here it «loes not deserve further conmuuents.
1994, Girardi Bertelli 1998 and references therein) and here it does not deserve further comments.
However. Fig.
However, Fig.
2 compares the| Juniuosities given iu Fig.l with similar results but by Girardi et al. (
2 compares the luminosities given in Fig.1 with similar results but by Girardi et al. (
1998: GOS hereinafter).
1998: G98 hereinafter).
To our surprise. oue finds that GOs Iuniuosiies appear svstemiatically züuter than in C99 bv about A og L/L. ~ 0.1 : 0.15. which is far from being a ποσο]le amount and. m turn. it appears of the right amount to solve the already quotec AMG? discrepancy.
To our surprise, one finds that G98 luminosities appear systematically fainter than in C99 by about $\Delta$ log $_{\odot}$ $\sim$ 0.1 $\div$ 0.15, which is far from being a negligible amount and, in turn, it appears of the right amount to solve the already quoted M67 discrepancy.
Such an evidence prompted us to investigate siniülar data in he literature. aiming to find the origin of such a differeice.
Such an evidence prompted us to investigate similar data in the literature, aiming to find the origin of such a difference.
To this purpose. the sane Fie.
To this purpose, the same Fig.
2 shows the
\ref{girardi} shows the
he DLAs in our sample have significant amount of dust. and have uigher abundances as compared to the rest of the sample.
the DLAs in our sample have significant amount of dust, and have higher abundances as compared to the rest of the sample.
The reason for the difference between the dust content of samples investigated by YOG and the DLA sample studied in his paper could possibly be the difference in the Mg II equivalent widths and redshifts of the two samples.
The reason for the difference between the dust content of samples investigated by Y06 and the DLA sample studied in this paper could possibly be the difference in the Mg II equivalent widths and redshifts of the two samples.
Mg Π A2796 is covered in only 22 of the SDSS spectra of the DLA sample. so we are unable to get meaningful value of the average equivalent width.
Mg II $\lambda2796$ is covered in only 22 of the SDSS spectra of the DLA sample, so we are unable to get meaningful value of the average equivalent width.
The average redshift of S3 is 2.84 (Table 13.
The average redshift of S3 is 2.84 (Table 1).
From Table | of YO6 we see that their full sample. sample |. has an average redshift of 1.33 and {οV) of 0.01.
From Table 1 of Y06 we see that their full sample, sample 1, has an average redshift of 1.33 and $E(B-V)$ of 0.01.
Assuming that the redshift dependence of absorber restframe (1513 derived by Ménnard et al. (
Assuming that the redshift dependence of absorber restframe $E(B-V)$ derived by Ménnard et al. (
2008) namely. 0V)ox(1|za)tt for zo«2 is valid at higher redshifts. sample $3 (assuming it has similar average Mg TT equivalent width) can be expected to have οV) of ~ 0.006 which is higher than the /Z(/5.—V) obtained for the full sample.
2008) namely, $E(B-V)\propto (1+{\rm z}_{ab})^{-1.1}$ for $_{ab}<2$ is valid at higher redshifts, sample S3 (assuming it has similar average Mg II equivalent width) can be expected to have $E(B-V)$ of $\sim$ 0.006 which is higher than the $E(B-V)$ obtained for the full sample.
As found here. most DLAs do not have significant amount of dust while of the observed DLAs do have some dust and likely higher abundances.
As found here, most DLAs do not have significant amount of dust while of the observed DLAs do have some dust and likely higher abundances.
Some of these could be the dusty systems found by Srianand et al. (
Some of these could be the dusty systems found by Srianand et al. (
2008).
2008).
It is possible that there is a bimodal distribution of dust in DLAs as was pointed out by Khare et al. (
It is possible that there is a bimodal distribution of dust in DLAs as was pointed out by Khare et al. (
2007): there may be systems with much higher dust content and QSOs behind such dusty systems are too faint to be seen in magnitude limited surveys and/or are too reddened to be included in the colour selected samples.
2007): there may be systems with much higher dust content and QSOs behind such dusty systems are too faint to be seen in magnitude limited surveys and/or are too reddened to be included in the colour selected samples.
Even though no evidence for this is found from studies of radio selected samples of QSOs (e.g. Ellison et al.
Even though no evidence for this is found from studies of radio selected samples of QSOs (e.g. Ellison et al.
2001: Akerman et al.
2001; Akerman et al.
2005: Jorgenson et al.
2005; Jorgenson et al.
2006). it may be worth noting that a few of such QSOs have been observed by SDSS through other means of selection (see e.g. Noterdaeme et al.20092) and have higher dust content and higher abundances.
2006), it may be worth noting that a few of such QSOs have been observed by SDSS through other means of selection (see e.g. Noterdaeme et al.2009a) and have higher dust content and higher abundances.
The SDSS data base would thus have very few of such systems.
The SDSS data base would thus have very few of such systems.
Ménnard et al. (
Ménnard et al. (
2008) have shown that the fraction of missed absorbers rises from | to with increase in Mg II rest equivalent width from | to 6A.
2008) have shown that the fraction of missed absorbers rises from 1 to with increase in Mg II rest equivalent width from 1 to 6.
We find a significantly higher dust content in DLAs along
We find a significantly higher dust content in DLAs along
a larger group, we build a unified model for the Mstars- Mingau relation, which describes the evolution of galaxy mass as a function of halo mass at any given redshift.
a larger group, we build a unified model for the $M_{stars}$ $M_{infall}$ relation, which describes the evolution of galaxy mass as a function of halo mass at any given redshift.
Satellite stellar mass is determined by the Mstars-Minfalt relation of central galaxy at the time of its infall.
Satellite stellar mass is determined by the $M_{stars}$ $M_{infall}$ relation of central galaxy at the time of its infall.
The best-fit model of both SDSS and VVDS stellar mass functions and clustering functions gives an obvious evolution of Mstars- Minfau relation from z=0.8 to z=0, with the mass of galaxies at higher redshift being lower than the galaxy mass at the present day, for all masses of hosting haloes.
The best-fit model of both SDSS and VVDS stellar mass functions and clustering functions gives an obvious evolution of $M_{stars}$ $M_{infall}$ relation from $z=0.8$ to $z= 0$, with the mass of galaxies at higher redshift being lower than the galaxy mass at the present day, for all masses of hosting haloes.
The Mstars-Minfau relation provided by the unified model is different as the relations in the models when SDSS and VVDS data are fitted separately, especially for galaxies with massive haloes and at high redshift.
The $M_{stars}$ $M_{infall}$ relation provided by the unified model is different as the relations in the models when SDSS and VVDS data are fitted separately, especially for galaxies with massive haloes and at high redshift.
The central galaxies are less massive and satellite galaxies are more massive in the unified model than the separate model.
The central galaxies are less massive and satellite galaxies are more massive in the unified model than the separate model.
Different sets of parameters can both give reasonable fits to the observed data, because the statistics of high mass galaxies are not tightly constrained in observation at high redshift.
Different sets of parameters can both give reasonable fits to the observed data, because the statistics of high mass galaxies are not tightly constrained in observation at high redshift.
We study the amount of galaxy stellar mass growth from z~0.8 to today, in either way of galaxy merger or star formation.
We study the amount of galaxy stellar mass growth from $z\sim0.8$ to today, in either way of galaxy merger or star formation.
For the models when SDSS and VVDS data are fitted separately, we find that for galaxies that reside in haloes of mass less than 10127!Mo, the masses of their most massive progenitors at z—0.83 vary from about 20 percent to 60 percent of the present day mass.
For the models when SDSS and VVDS data are fitted separately, we find that for galaxies that reside in haloes of mass less than $10^{12}h^{-1}M_{\odot}$, the masses of their most massive progenitors at $z=0.83$ vary from about 20 percent to 60 percent of the present day mass.
Meanwhile, galaxy mergers contribute only a small fraction to the galaxy mass of today.
Meanwhile, galaxy mergers contribute only a small fraction to the galaxy mass of today.
This indicates that a large fraction of z=0 stellar masses comes from star formation during the period between z—0.83 and z—0.
This indicates that a large fraction of $z=0$ stellar masses comes from star formation during the period between $z=0.83$ and $z=0$.
For galaxies within massive haloes, the total amount of stellar mass from the main progenitor at z—0.83 and from that increased from mergers with other galaxies actually exceeds the present-day galaxy mass.
For galaxies within massive haloes, the total amount of stellar mass from the main progenitor at $z=0.83$ and from that increased from mergers with other galaxies actually exceeds the present-day galaxy mass.
Although there could be a difference of the definitions for stellar mass of central galaxies in observation and in the models based on simulation (for example, the stars in the envelope of the central galaxies may not be counted in observation), this indicates that there may be little star formation ongoing in these massive galaxies.
Although there could be a difference of the definitions for stellar mass of central galaxies in observation and in the models based on simulation (for example, the stars in the envelope of the central galaxies may not be counted in observation), this indicates that there may be little star formation ongoing in these massive galaxies.
The unified model basically tells the same story, except that the contribution of the most massive progenitors at z—0.83
The unified model basically tells the same story, except that the contribution of the most massive progenitors at $z=0.83$
arises from the inclusion of the UV Ilux here).
arises from the inclusion of the UV flux here).
This value again suggests a U Gem classification for οV1.
This value again suggests a U Gem classification for CV1.
The orbital period distribution of CVs typically ranges from tens of minutes up to about 15 bh. Phe value of £46 fora CV can reveal much about the parameters of thesystem. (e.g.Warner 1995).
The orbital period distribution of CVs typically ranges from tens of minutes up to about 15 h. The value of $P_{\mathrm{orb}}$ for a CV can reveal much about the parameters of thesystem \citep[e.g.][]{warner95}.
. Whilst it is challenging to cürectly measure the orbital period for a svstem as faint and. crowded as CVI in M22. we can none the less attempt to constrain it bv comparing the optical. UV and X-ray. properties of the svstem to those of Galactic CVs with known orbital periods.
Whilst it is challenging to directly measure the orbital period for a system as faint and crowded as CV1 in M22, we can none the less attempt to constrain it by comparing the optical, UV and X-ray properties of the system to those of Galactic CVs with known orbital periods.
For example. the absolute magnitude at minimum of the source. Aly~5.4 mag (Ancersonctal. 2003).. suggests an orbital period of Pan=1343 h (see equation 18 of Warner (1987))).
For example, the absolute magnitude at minimum of the source, $M_{\mathrm{V}} \sim 5.4$ mag \citep{ack03}, , suggests an orbital period of $P_{\mathrm{orb}} = 13\pm3$ h (see equation 18 of \citet{warner87}) ).
Furthermore. because of the increased UV Luminosity of CVs at longer orbital periods due to their corresponcdinglv higher acerction rates (vanTeeseling.etal.1996)... and the lack of any clear correlation between X-ray. luminosity and orbital period. it would. seem reasonableto expect a correlation between ἐς ffiopr and orbital period for non-magnetic CVs.
Furthermore, because of the increased UV luminosity of CVs at longer orbital periods due to their correspondingly higher accretion rates \citep{teese96}, and the lack of any clear correlation between X-ray luminosity and orbital period, it would seem reasonableto expect a correlation between $ F_{\mathrm{X}} $ $ F_{\mathrm{uv + opt}}$ and orbital period for non-magnetic CVs.
In fact. vanTeeselingctal. have. found evidence for such a correlation: their fig.
In fact, \citeauthor{teese96} have found evidence for such a correlation: their fig.
6 shows that for a sample of non-magnetic field CVs. with increasing orbital period there is a trend of decreasing fs ffi a.
6 shows that for a sample of non-magnetic field CVs, with increasing orbital period there is a trend of decreasing $F_{\mathrm{X}}$ $F_{\mathrm{uv + opt}}$ .
Interestingly however. our PN/Popt value for CVI places it in the region of the vanTeeselingetal. plot with usZO2 he in contradiction to our previous orbital-period constraints.
Interestingly however, our $F_{\mathrm{X}}$ $F_{\mathrm{uv + opt}}$ value for CV1 places it in the region of the \citeauthor{teese96} plot with $P_{\mathrm{orb}}\lesssim 2$ h, in contradiction to our previous orbital-period constraints.
We discuss below two wavs in which this could be resolved: ]lence. we conclude that our classification of CVI as a dwarf nova of U Cem type remains secure.
We discuss below two ways in which this could be resolved: Hence, we conclude that our classification of CV1 as a dwarf nova of U Gem type remains secure.
Strong. A4686 line enmüssion is characteristic of magnetic CVs (inpolars.itiscomparableinstrength.to11:32:Warner 1995).
Strong $\lambda 4686$ line emission is characteristic of magnetic CVs \citep[in polars, it is comparable in strength to H$\beta$;][]{warner95}.
. In LPs. for example. this high-excitation line is most likely powered by X-ray heatingof the regions of the aceretion curtains close to the magnetic poles. interior to the inner edge of the magnetically truncated accretion disc (Saitoetal.2010).
In IPs, for example, this high-excitation line is most likely powered by X-ray heatingof the regions of the accretion curtains close to the magnetic poles, interior to the inner edge of the magnetically truncated accretion disc \citep{saito10}.
. The strongHer. Ho and. Balmer emission (see Figs.
The strong, $\alpha{}$ and Balmer emission (see Figs.
S and 10)) of VIOL are characteristic of CVs in quiescence. while the lack of any obvious emission an upper limit of G can be placed. on the equivalent width of any A4686 cmiuission in the spectrum sugeests that the system. is unlikely το be maenetic.
\ref{fig:redspec} and \ref{fig:bluespec}) ) of V101 are characteristic of CVs in quiescence, while the lack of any obvious emission – an upper limit of $6$ can be placed on the equivalent width of any $\lambda 4686$ emission in the spectrum – suggests that the system is unlikely to be magnetic.
The enhanced contribution of the aceretion disc to the Iuminosity of the system at shorter wavelengths means we would expect the absorption lines of the secondary star to be more apparent at longer wavelengths.
The enhanced contribution of the accretion disc to the luminosity of the system at shorter wavelengths means we would expect the absorption lines of the secondary star to be more apparent at longer wavelengths.
Ehe only possible absorption feature from the secondary which can be identified in our summed spectrum is a weak feature ab TLISA although we cannot be certain this is not an artefact.
The only possible absorption feature from the secondary which can be identified in our summed spectrum is a weak feature at 7118, although we cannot be certain this is not an artefact.
Our phased. optical light curves in the ancl à bands show varying cleerces of modulation at the orbital period of Neillctal.(2002).
Our phased optical light curves in the and bands show varying degrees of modulation at the orbital period of \cite{neill02}.
. Phe main peak in the light curve may be due to the hotspot (seec.g.Woodetal.1986).
The main peak in the light curve may be due to the hotspot \citep[see e.g.][]{wood86}.
. Comparing our PN £Poptc0.06 to the trend in fig.
Comparing our $F_{\mathrm{X}}$ $F_{\mathrm{uv + opt}}\geq0.06$ to the trend in fig.
6 of vanTeeselingetal.(1996). as we did for M22 €CVI. again we find that for our result. their trend would predict a shorter orbital period ol up to a maximum of around three hours. compared to the measured 2,4,=5.796x0.036 h (Neilletal.2002).
6 of \citet{teese96}, as we did for M22 CV1, again we find that for our result, their trend would predict a shorter orbital period of up to a maximum of around three hours, compared to the measured $P_{\mathrm{orb}} = 5.796\pm0.036$ h \citep{neill02}.
. This discrepancy is not as pronounced as what we found for the M22 CY. although in that case we did not have a measured has το compare with.
This discrepancy is not as pronounced as what we found for the M22 CV, although in that case we did not have a measured $P_{\mathrm{orb}}$ to compare with.
However. again we note the lack of U CGem-type systems in the vanTeeselingetal.(1996). sample.
However, again we note the lack of U Gem-type systems in the \citet{teese96} sample.
For M22 CVI. the X-ray. properties confirm a CV rather han an LAINB nature.
For M22 CV1, the X-ray properties confirm a CV rather than an LMXB nature.
The high Lx and X-ray spectra iwdness are indicative of either a quiescent DN or an LDP.
The high $L_{\mathrm{X}}$ and X-ray spectral hardness are indicative of either a quiescent DN or an IP.
The outburst characteristics are consistent with a norma ong Paw Cz10 h) dwarf nova of U Gem type. but canno o very easily reconciled with any particular example of the known IPs.
The outburst characteristics are consistent with a normal long $P_{\mathrm{orb}}$ $\gtrsim 10$ h) dwarf nova of U Gem type, but cannot be very easily reconciled with any particular example of the known IPs.
We find the expected. UV. excess compared. to he cluster main sequence in 335 versus (€y36 Voss). ane hat the (€ 13 colour is consistent with a longer orbita »eriod.
We find the expected UV excess compared to the cluster main sequence in $U_{336}$ versus $U_{336}-V_{555}$ ), and that the $U-V$ ) colour is consistent with a longer orbital period.
For the M5 CW. VIOL. our optical spectra are tvpica or a quiescent dwarf nova.
For the M5 CV, V101, our optical spectra are typical for a quiescent dwarf nova.
The lack of any significan emission suggests the svstem is not magnetic.
The lack of any significant emission suggests the system is not magnetic.
The modulation we observe in the #-bancl light) curve is consistent with the orbital period of 5.796x0.036 hi. Loune by Neilletal.(2002).
The modulation we observe in the -band light curve is consistent with the orbital period of $5.796\pm0.036$ h found by \citet{neill02}.
. As this remains the elobular cluster CV most amenable to detailed. study. from the &eround in quiescence. it will benefit. from future higherS/N. optical spectroscopy.
As this remains the globular cluster CV most amenable to detailed study from the ground in quiescence, it will benefit from future higherS/N phase-resolved optical spectroscopy.
We wouldlike to thank the referee. ChristianIxnigge.[or helpful comments and suggestions.
We wouldlike to thank the referee, ChristianKnigge,for helpful comments and suggestions.
This research has made use of data obtained from the and software provided by the Center (CXC) in the application package CLAO.
This research has made use of data obtained from the and software provided by the ) in the application package .
This research has also mace
This research has also made
Another eritical element for all possible explanations is hydrogen.
Another critical element for all possible explanations is hydrogen.
We note that in all objects the abundance of hydrogen relative to the metals is less than the solar value. even in the one case where it is detected.
We note that in all objects the abundance of hydrogen relative to the metals is less than the solar value, even in the one case where it is detected.
This ts consistent with accretion. of predominantly volatile-depleted material. since hydrogen accreted from the interstellar matter or from circumstellar water/ice would stay in the convection zone and can only accumulate with time.
This is consistent with accretion of predominantly volatile-depleted material, since hydrogen accreted from the interstellar matter or from circumstellar water/ice would stay in the convection zone and can only accumulate with time.
In the case of the DBZ star GD40. Kleinetal.(2010) were able to deduce very detailed conclusions from à comparison of photospherie abundances with those of various solar system bodies.
In the case of the DBZ star GD40, \cite{Klein.Jura.ea10} were able to deduce very detailed conclusions from a comparison of photospheric abundances with those of various solar system bodies.
The abundance ratios we obtain here roughly follow that of the bulk Earth.
The abundance ratios we obtain here roughly follow that of the bulk Earth.
However. the remaining uncertainties of surface gravity. abundances. and time since the last aceretion event do not allow us to draw any further conclusions in the present study.
However, the remaining uncertainties of surface gravity, abundances, and time since the last accretion event do not allow us to draw any further conclusions in the present study.
High resolution. high S/N observations and extending the spectral coverage at least to 3200A.. or better to 2700 ((to cover the resonance lines) are needed to refine the element abundances and detect the very crucial element silicon. and if possible carbon and oxygen.
High resolution, high S/N observations and extending the spectral coverage at least to 3200, or better to 2700 (to cover the resonance lines) are needed to refine the element abundances and detect the very crucial element silicon, and if possible carbon and oxygen.
We can. however. estimate a lower limit for the total accreted mass based on the available data for the convection zone.
We can, however, estimate a lower limit for the total accreted mass based on the available data for the convection zone.
For SDSS0956+5912. which has the highest metal pollution. we find a total mass for theobserved metals Μα. Ca. Fe. Na of 1.5 I0?g. Adding oxygen with the same ratio to Mg as in GD4O (Kleinetal.2010) this number becomes 4.8 107g. even more than in the extremely heavily polluted SDSS$07384+1835 (Dufouretal.2010).
For SDSS0956+5912, which has the highest metal pollution, we find a total mass for the metals Mg, Ca, Fe, Na of $1.5\,10^{23}$ g. Adding oxygen with the same ratio to Mg as in GD40 \citep{Klein.Jura.ea10} this number becomes $4.8\,10^{23}$ g, even more than in the extremely heavily polluted SDSS0738+1835 \citep{Dufour.Kilic.ea10}.
. The total amount of H in the convection zone is 1.0 107g. For SDSS114441218. which has the lowest abundances. these numbers are 3.2 107g (observed metals). and 7.9 10e (including oxygen).
The total amount of H in the convection zone is $1.0\,10^{24}$ g. For SDSS1144+1218, which has the lowest abundances, these numbers are $3.2\,10^{20}$ g (observed metals), and $7.9\,10^{20}$ g (including oxygen).
These masses span the range of the most massive asteroids in our own planetary system.
These masses span the range of the most massive asteroids in our own planetary system.
These are the absolute minimum of the accreted masses — depending on how long ago the aceretion ended. these masses can be at least two orders of magnitude larger. which will require minor planets much larger than known in the Solar system.
These are the absolute minimum of the accreted masses – depending on how long ago the accretion ended, these masses can be at least two orders of magnitude larger, which will require minor planets much larger than known in the Solar system.
Another open question. recently discussed by Farihietal.(2011) in the context of G77-50. a DAZ of similar low temperature. is how such à very massive asteroid is suddenly driven into its host star from an orbit apparently stable over the past 5 Gyrs.
Another open question, recently discussed by \cite{Farihi.Dufour.ea11} in the context of G77-50, a DAZ of similar low temperature, is how such a very massive asteroid is suddenly driven into its host star from an orbit apparently stable over the past 5 Gyrs.
ou the Canada-France-Tawaii Telescope use a curvature sensor to measure wavefrout error and have traclitionally used a bimorph mirror with a relatively simall umber of degrees of freedom (50).
on the Canada-France-Hawaii Telescope use a curvature sensor to measure wavefront error and have traditionally used a bimorph mirror with a relatively small number of degrees of freedom $<$ 50).
For several reasons mclucdiug the use of avadancl photodiodes aud the need to divide the helt iuto fewer elements. curvature based svstenmis can operate on fainter targets and have historically been more successtul on extragalactic targets.
For several reasons including the use of avalanch photodiodes and the need to divide the light into fewer elements, curvature based systems can operate on fainter targets and have historically been more successful on extragalactic targets.
Iu any AO system. only a fraction of the light is restored to the diffraction-linuted core. while the rest remains iu an extended halo roughly the size of the seeing disk.
In any AO system, only a fraction of the light is restored to the diffraction-limited core, while the rest remains in an extended halo roughly the size of the seeing disk.
A paraiueter used to describe the effectiveuess of AO is the Strohl ratio.
A parameter used to describe the effectiveness of AO is the Strehl ratio.
It is the ratio of the peal intensity of the poiut-spread function (PSF) compared to the theoretical peak intensity of a diffraction-linited source.
It is the ratio of the peak intensity of the point-spread function (PSF) compared to the theoretical peak intensity of a diffraction-limited source.
A hielo:order system with many actuators (or degrees of freedoni) will normally result iu à higher Strelil ratio than a lower-order system on the same telescope.
A higher-order system with many actuators (or degrees of freedom) will normally result in a higher Strehl ratio than a lower-order system on the same telescope.
Typical Strehl ratios at 1.6 son for the Neck system are around 0.3 for the guide star and fall off with angular separation.
Typical Strehl ratios at 1.6 $\mu$ m for the Keck system are around 0.3 for the guide star and fall off with angular separation.
The Strehl ratio is typically below 0.1 at a distance of 30" from the euice star. although there is significant variation with atmospheric conditions.
The Strehl ratio is typically below 0.1 at a distance of $30''$ from the guide star, although there is significant variation with atmospheric conditions.
The Ikeck system requires a guide star brighter than 12th mae at AR for full correction.
The Keck system requires a guide star brighter than 12th mag at $R$ for full correction.
Partial correction is possible on guide stars as faint as Lith mag.
Partial correction is possible on guide stars as faint as 14th mag.
This iuitation of using bright sources for wavefrout references its severely dinüted the use of ΑΟ svstenis on faint extragalactic targets;
This limitation of using bright sources for wavefront references has severely limited the use of AO systems on faint extragalactic targets.
Our study of faint field ealaxics is he first of its kind aud is only possible because of the eh density of ealaxies on the slaw.
Our study of faint field galaxies is the first of its kind and is only possible because of the high density of galaxies on the sky.