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Therefore. conduction appears to be more important than. for example. mixing. which would also smooth the metalicity jump. | Therefore, conduction appears to be more important than, for example, mixing, which would also smooth the metalicity jump. |
In numerical simulations of cold fronts. often more than one cold front is seen on both sides of the cluster (e.g.?).. | In numerical simulations of cold fronts, often more than one cold front is seen on both sides of the cluster \citep[e.g.][]{tittley2005}. |
If there are more then one. they are expected to alternate on an axis through the centre of the cluster. | If there are more then one, they are expected to alternate on an axis through the centre of the cluster. |
In Abell 2052. the counterpart of the front at 3.2 may lie North-East of the centre. | In Abell 2052, the counterpart of the front at $^{\prime}$ may lie North-East of the centre. |
? identified two jumps in surface brightness in that area using a deep Chandra image. | \citet{blanton2009}
identified two jumps in surface brightness in that area using a deep Chandra image. |
The locations of these jumps are indicated in Fig. 9.. | The locations of these jumps are indicated in Fig. \ref{fig:acis}. |
They are located at 45" and 67” from the cluster centre. | They are located at $^{\prime\prime}$ and $^{\prime\prime}$ from the cluster centre. |
The 67" jump may be the counterpart of the cold front at 3.2’. | The $^{\prime\prime}$ jump may be the counterpart of the cold front at $^{\prime}$. |
If we look back to Fig. 3.. | If we look back to Fig. \ref{fig:profiles}, |
then the iron abundance profile of the NE side of the central galaxy shows ἃ jump around 1.7'. | then the iron abundance profile of the NE side of the central galaxy shows a jump around $^{\prime}$. |
Probably. this jump ts a bit too far from the jumps identified in the Chandra data to be associated with the counterpart of the cold front we find. although the spatial resolution of our iron abundance profile is relatively low. | Probably, this jump is a bit too far from the jumps identified in the Chandra data to be associated with the counterpart of the cold front we find, although the spatial resolution of our iron abundance profile is relatively low. |
In addition. this North-Eastern region may be affected by the AGN activity in the core of the cluster. | In addition, this North-Eastern region may be affected by the AGN activity in the core of the cluster. |
However. the position of the NE jumps in surface brightness and iron abundance appear to support the core oscillation interpretation. | However, the position of the NE jumps in surface brightness and iron abundance appear to support the core oscillation interpretation. |
cloud. — which. as discussed in refnocomb.. is greater than 1: it is 3.5 c 1. as predicted by the ambipolar-cdillusion theory. | cloud – which, as discussed in \\ref{nocomb}, is greater than 1; it is 3.5 $\pm$ 1, as predicted by the ambipolar-diffusion theory. |
We have shown/explained that: We are grateful to Nicholas Chapman for providing Figures 1. and 2 and for invaluable discussions on observational issues. | We have shown/explained that: We are grateful to Nicholas Chapman for providing Figures 1 and 2 and for invaluable discussions on observational issues. |
TMS work was supported in part by the National Science Foundation under grant. NSE AST-07-09206 to the University. of Hlinois. | TM's work was supported in part by the National Science Foundation under grant NSF AST-07-09206 to the University of Illinois. |
Part of this work was carried. out at the Jet Propulsion Laboratory. California Institute of ‘Technology. under a contract with the National Acronautics and Space Administration. | Part of this work was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration. |
derived in this way. | derived in this way. |
For the rest of the stars we used the (4—L’), vs. Tey calibration also described by Ohnaka&Tsuji(1996). | For the rest of the stars we used the $(J-L')_o$ vs. $_{\rm{eff}}$ calibration also described by \citet{ohn96}. |
. Infrared. photometry for our stars is taken [rom Noguchietal.(1981). and Fouquéetal.(1992). | Infrared photometry for our stars is taken from \citet{nog81} and \citet{fou92}. |
. Effective temperatures for J-stars derived in (his wav do not dilfer significantly from those derived in N- aud SC-tvpe carbon stars. | Effective temperatures for J-stars derived in this way do not differ significantly from those derived in N- and SC-type carbon stars. |
The estimated error in Γι is £150 Ix (see Olnaka Tsuji 1996. for details), | The estimated error in $_{\rm{eff}}$ is $\pm 150$ K (see Ohnaka Tsuji 1996, for details). |
The set of models used in (his analysis was computed by the Uppsala group (see [Eriksson et al. | The set of models used in this analysis was computed by the Uppsala group (see Eriksson et al. |
1984. for details). | 1984, for details). |
The models cover the Teg=2500—3500 Ix. C/O=1.0—1.35 ranges and all have the same gravity log ο--0.0. | The models cover the $_{\rm{eff}}=2500-3500$ K, $=1.0-1.35$ ranges and all have the same gravity log $=0.0$. |
The input elemental abundances adopted for the J-star models were (he solar values. with the exception of C. N and O which were assumed to be altered relative to the Sun. | The input elemental abundances adopted for the J-star models were the solar values, with the exception of C, N and O which were assumed to be altered relative to the Sun. |
The CNO abundances in the model atinosphere for a given star were taken from the literature (Lambertοἱal.1986:Abia&Isern1997). | The CNO abundances in the model atmosphere for a given star were taken from the literature \citep{lam86,abi97}. |
. For each star a model atmosphere was interpolated in T and C/O ratio in this erid. | For each star a model atmosphere was interpolated in $_{\rm{eff}}$ and C/O ratio in this grid. |
A ivpical mieroturbulence velocity for AGB stars €=3 | was adopted or taken from the literature when available (Lambertetal.10560). | A typical microturbulence velocity for AGB stars $\xi=3$ $^{-1}$ was adopted or taken from the literature when available \citep{lam86}. |
. Table 2 shows the atmosphere parameters used in the analvsis. | Table 2 shows the atmosphere parameters used in the analysis. |
The greatest difficulty in the analvsis of atomic lines of carbon stars lies in the heavy blend effect of molecular bands. | The greatest difficulty in the analysis of atomic lines of carbon stars lies in the heavy blend effect of molecular bands. |
The problem of blending is even more serious in J-stars because of the strong Cy and CN band absorptions. | The problem of blending is even more serious in J-stars because of the strong $_2$ and CN band absorptions. |
For Chis reason. identification of spectral lines in carbon stars is very difficult ancl. moreover. detailed laboratory data of these molecules are usually scarce. | For this reason, identification of spectral lines in carbon stars is very difficult and, moreover, detailed laboratory data of these molecules are usually scarce. |
and temporal distributions of supernova events in low and high mass resolution simulations will produce structurally different galaxies. | and temporal distributions of supernova events in low and high mass resolution simulations will produce structurally different galaxies. |
Furthermore. the high spatial detail produced by using large numbers of particles may allow for the formation of non-spherically svaunetric structures. such as galactic winds. | Furthermore, the high spatial detail produced by using large numbers of particles may allow for the formation of non-spherically symmetric structures, such as galactic winds. |
Such structures would likewise be affected by resolution induced changes to the artificial viscosity of gas particles (2) and changes to (he disk scale jeight. and the geometry of the disk (?7).. | Such structures would likewise be affected by resolution induced changes to the artificial viscosity of gas particles \citep{Kaufmann07} and changes to the disk scale height and the geometry of the disk \citep{MacLow88, MacLow89}. |
In (his paper. we explore (he effects of force and mass resolution on S06 SF aud stellar eedback in different mass galaxies. | In this paper, we explore the effects of force and mass resolution on S06 SF and stellar feedback in different mass galaxies. |
When analvzing galaxies wilh (his SF recipe. il is recessary (o know the lower limit of mass ancl force resolution above which SF and the stellar disk properties converge. | When analyzing galaxies with this SF recipe, it is necessary to know the lower limit of mass and force resolution above which SF and the stellar disk properties converge. |
Correctly determining the amount and location of SF is vital in simulations in galaxies because SF 1) produces and distributes metals throughout the galaxy. 2) affects the distribution of matter in the galaxy through feedback. and 3) enables us to relate simulations to observations. | Correctly determining the amount and location of SF is vital in simulations in galaxies because SF 1) produces and distributes metals throughout the galaxy, 2) affects the distribution of matter in the galaxy through feedback, and 3) enables us to relate simulations to observations. |
Therelore. we ask: what resolutions are sufficient to establish the galactic SF rates. the history of SF. the amount of stellar feedback. and the distribution ol stellar and gaseous matter? | Therefore, we ask: what resolutions are sufficient to establish the galactic SF rates, the history of SF, the amount of stellar feedback, and the distribution of stellar and gaseous matter? |
To this end. we simulated. a series of isolated galaxies of different. masses at different. mass ancl force resolutions ancl analvzed properties relating to SE. | To this end, we simulated a series of isolated galaxies of different masses at different mass and force resolutions and analyzed properties relating to SF. |
Isolated galaxies have the advantage of requiring less computational expense while allowing us lo separate out environmental effects from those caused by different resolutions. | Isolated galaxies have the advantage of requiring less computational expense while allowing us to separate out environmental effects from those caused by different resolutions. |
We then analvze the SF and stellar feedback in the models and relate them to the history and structure of the galaxies to determine the resolution necessary [or convergence and to describe the effects of resolution. | We then analyze the SF and stellar feedback in the models and relate them to the history and structure of the galaxies to determine the resolution necessary for convergence and to describe the effects of resolution. |
The outline of our paper is as follows. | The outline of our paper is as follows. |
In 82. we introduce the models aud describe the computational methods used. | In 2, we introduce the models and describe the computational methods used. |
The results of these simulations are described in terms of global SF (83.1).80.1, stellar feedback (83.2)(SOS, and stellar distribution (83.3). | The results of these simulations are described in terms of global SF 3.1), stellar feedback 3.2) and stellar distribution 3.3). |
(39.9, We address the phlivsical connection between the SF. feedback. and structure and discuss implications of these results for large volume simulations in 84. | We address the physical connection between the SF, feedback, and structure and discuss implications of these results for large volume simulations in 4. |
We conclude 85 with a list of recommendations for future cosmological simulations. | We conclude 5 with a list of recommendations for future cosmological simulations. |
Our set of isolated galaxies ranges in mass from 10. to LOMAL. | Our set of isolated galaxies ranges in mass from $10^9 M_\odot$ to $10^{13} M_\odot$. |
We simulated each galaxy ab five mass resolutions with a range of 50 to 10° DM particles and with the same number of initial gas particles. | We simulated each galaxy at five mass resolutions with a range of 50 to $10^5$ DM particles and with the same number of initial gas particles. |
For simplicitw. we reler to the mass resolution of a given simulation bv the number of DM particles with the understanding that this also denotes | For simplicity, we refer to the mass resolution of a given simulation by the number of DM particles with the understanding that this also denotes |
were plotted on a V vs V—7 CM diagram. | were plotted on a $V$ vs $V-I$ CM diagram. |
All but seven of the optically detected stars are foreground stars. | All but seven of the optically detected stars are foreground stars. |
Since seven sources is not statistically significant [or (he voung cluster. we instead plot the J vs J—HCM diagram in Fie. | Since seven sources is not statistically significant for the young cluster, we instead plot the $J$ vs $J-H$ CM diagram in Fig. |
6. | 6. |
In (his plot the slanting arrow denotes the reddening vector associated with 15 mae of extinction. | In this plot the slanting arrow denotes the reddening vector associated with 15 mag of extinction. |
YSOs identified in the Jiffy CC diagram in Fig.44 are shown as open circles: triangles mark the brieht stars. | YSOs identified in the $JHK$ CC diagram in 4 are shown as open circles; triangles mark the bright stars. |
PAIS evolutionary tracks (Palla&Stahler1999). [or ages 1. 3 and 5 Alves are shown in Fig. | PMS evolutionary tracks \citep{palla99} for ages 1, 3 and 5 Myrs are shown in Fig. |
6 with dotted. short-dashed. and lone-dashecl lines. respectively. | 6 with dotted, short-dashed, and long-dashed lines, respectively. |
We have assunied a distance of 4.2 kpe (Molinarietal.1996). and have reddened the isochrones with the average extinction of sly = 7.6 mag estimated [rom the Jif CC diagram analvsis in Sect. | We have assumed a distance of 4.2 kpc \citep{mol96} and have reddened the isochrones with the average extinction of $A_{\rm V}$ = 7.6 mag estimated from the $JHK$ CC diagram analysis in Sect. |
3.1. | 3.1. |
Given the poor statisües for this distant. cluster. it is not easy (o derive the age bv isochrone filling. | Given the poor statistics for this distant cluster, it is not easy to derive the age by isochrone fitting. |
Even for regions with good statistics. it is «uite difficult to constrain the age because most low mass stars spend the bulk of their PAIS Gime on the Iavashi track. which is thevertieal part of each isochrone. | Even for regions with good statistics, it is quite difficult to constrain the age because most low mass stars spend the bulk of their PMS time on the Hayashi track, which is the part of each isochrone. |
Llowever. a small fraction of observed voung stars will coincide with the brief evolutionary plase associated wilh the IIenvey track. which is thehorizontal transition between the Iavashi track and the Zero-Age-Main Sequence (ZANIS) (see e.g, Ascensoetal. 2007)). | However, a small fraction of observed young stars will coincide with the brief evolutionary phase associated with the Henyey track, which is the transition between the Hayashi track and the Zero-Age-Main Sequence (ZAMS) (see e.g. \citealt*{ascenso07}) ). |
In Fig. | In Fig. |
6 we note that there is a large population of low mass/low J magnitude candidate voung stus (open circles) that lie to the of the 1. Myr isochrone (dotted line). | 6 we note that there is a large population of low mass/low $J$ magnitude candidate young stars (open circles) that lie to the of the 1 Myr isochrone (dotted line). |
Some of these sources max be Class I protostars associated with the cluster. | Some of these sources may be Class I protostars associated with the cluster. |
However. most will be Class II sources (T Tauri stars) with an age of about 1 Myr: many of these sources lie to the right of the 1 Myr isochrone because of extinction which. to these embedded sources. will be higher than the A ~ 8 mags used to deredden the PAIS isochrones. | However, most will be Class II sources (T Tauri stars) with an age of about 1 Myr; many of these sources lie to the right of the 1 Myr isochrone because of extinction which, to these embedded sources, will be higher than the $A_{\rm V}$ $\sim$ 8 mags used to deredden the PMS isochrones. |
However. there is also a population of candidate low mass voung stus that lie to the/eff of the 1 Myr isochrone. | However, there is also a population of candidate low mass young stars that lie to the of the 1 Myr isochrone. |
These sources represent a more evolved group of voung stars that are probably 3 Myr or older. | These sources represent a more evolved group of young stars that are probably 3 Myr or older. |
Indeed. there are five probable cluster members that align horizontally ancl mav therefore represent the Henvev part of the 3 Myr isochrone (short dash). | Indeed, there are five probable cluster members that align horizontally and may therefore represent the Henyey part of the 3 Myr isochrone (short dash). |
Hence. we estimate that there is à low mass population of stars associated with IRAS 193434-2026 that is best represented by an age of 1 Myr or more. | Hence, we estimate that there is a low mass population of stars associated with IRAS 19343+2026 that is best represented by an age of 1 Myr or more. |
The mass range plotted in the figure is [rom 0.6 to 3.0 lor the 3 Myr isochrone. | The mass range plotted in the figure is from 0.6 to 3.0 for the 3 Myr isochrone. |
The 5 Myr isochrone (long dash: mass range 0.6 2.5 )) may well be within the age spread of the cluster. | The 5 Myr isochrone (long dash; mass range 0.6 – 2.5 ) may well be within the age spread of the cluster. |
However. it extends bevond (he limits of the observed data points. | However, it extends beyond the limits of the observed data points. |
One should also note that the / vs J—II diagram contains only DODey sources that are common to the J ancl {1 bands. | One should also note that the $J$ vs $J-H$ diagram contains only 333 sources that are common to the $J$ and $H$ bands. |
This is a considerable under-representation of the full sample. | This is a considerable under-representation of the full sample, |
The first iron Ίνα lines were discovered in NGC 1151. and a few source with larec absorbing columus. iun which the line was thought to originate ((Mushotzky. Πο Serlemitsos 1978: Mushotzky 1982). | The first iron $\alpha$ lines were discovered in NGC 4151, and a few source with large absorbing columns, in which the line was thought to originate ((Mushotzky, Holt Serlemitsos 1978; Mushotzky 1982). |
The first unobscured ACN to show line emissiou was MCG-6-30-15 CNandra ct al. | The first unobscured AGN to show line emission was MCG-6-30-15 (Nandra et al. |
1989: Matsuoka et al. | 1989; Matsuoka et al. |
1990) and ssubsequeutlv found iron Ίνα emission to be extremely commen in Sevfert galaxies (Pounds et al. | 1990) and subsequently found iron $\alpha$ emission to be extremely common in Seyfert galaxies (Pounds et al. |
1990: Nandra Pounds 1991). | 1990; Nandra Pounds 1994). |
Line cuntission had been predicted frou optically-thick material close to the nucleus (Guilbert Rees 1988). including the accretion disk (Fabian ct al. | Line emission had been predicted from optically-thick material close to the nucleus (Guilbert Rees 1988), including the accretion disk (Fabian et al. |
1989). | 1989). |
Detailed predictions of the line strenetl from the disk (Georee Fabian 1991: Matt. Perola Piro 1991) were found to be in excellent agrecinent with the observations (Nandra Pounds 1991). but the ddata were unable to deteriiue the width or profile of these lines. | Detailed predictions of the line strength from the disk (George Fabian 1991; Matt, Perola Piro 1991) were found to be in excellent agreement with the observations (Nandra Pounds 1994), but the data were unable to determine the width or profile of these lines. |
This is of clear iuportanuce. as the profiles allow the location and geometry of the material to be constrained. | This is of clear importance, as the profiles allow the location and geometry of the material to be constrained. |
Specifically. in the case of au accretion disk. large widths and distinctive profiles are expected due to the rotation aud gravitational effects of the black hole (Fabian et al. | Specifically, in the case of an accretion disk, large widths and distinctive profiles are expected due to the rotation and gravitational effects of the black hole (Fabian et al. |
1989: Stella 1990: Laor 1991: Matt ct al. | 1989; Stella 1990; Laor 1991; Matt et al. |
1992). | 1992). |
The launch of ooffered an opportunity to test these models. with the SIS detectors having good | The launch of offered an opportunity to test these models, with the SIS detectors having good |
than the continuum because of the larger size of the emitting region, | than the continuum because of the larger size of the emitting region. |
As areference model. we use a singular tsothermal sphere with external shear (SIST). | As a reference model, we use a singular isothermal sphere with external shear (SIST). |
This is probably the simplest model capable to reproduce the observed positions and the flux ratio RR98). | This is probably the simplest model capable to reproduce the observed positions and the flux ratio R98). |
The parameters of this model can be found in Table 2.. | The parameters of this model can be found in Table \ref{tab:modpar}. |
The observational uncertainties lead to an internal error of only To examine the much larger possible errors due to the modeling. we used a more general approach of models consisting of a singular isothermal ellipsoidal mass distribution (SIEMD. see Kassiola Kovner 1993)) with external shear. | The observational uncertainties lead to an internal error of only To examine the much larger possible errors due to the modeling, we used a more general approach of models consisting of a singular isothermal ellipsoidal mass distribution (SIEMD, see Kassiola Kovner \cite{kassiola93}) ) with external shear. |
For the model-fitting we fixed ellipticity £ and shear y. and used the other parameters listed in Table 2 (including. the position angles) and the source position to fit the observations. | For the model-fitting we fixed ellipticity $\epsilon$ and shear $\gamma$, and used the other parameters listed in Table \ref{tab:modpar} (including the position angles) and the source position to fit the observations. |
Due to the small number of constraints. the position of the lensing galaxy was fixed at the observed values. | Due to the small number of constraints, the position of the lensing galaxy was fixed at the observed values. |
This was carried out for a range of values for € and y. | This was carried out for a range of values for $\epsilon$ and $\gamma$. |
With the restriction of£«0.3 and y«0.2. we find a maximum deviation of |20% and 10% for the time delay. | With the restriction of $\epsilon<0.3$ and $\gamma<0.2$, we find a maximum deviation of $+20\,\%$ and $-10\,\%$ for the time delay. |
As an example. the model with zero external shear is also given in Table 2.. | As an example, the model with zero external shear is also given in Table \ref{tab:modpar}. |
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