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Abadietal.(2003a.b) present detailed analyses of simulated galaxies with kinematic and photometric properties similar to observed Sab galaxies.
\cite{abadi03a, abadi03b} present detailed analyses of simulated galaxies with kinematic and photometric properties similar to observed Sab galaxies.
Although the various attempts to simulate disk formation have provided some impressive successes. the essential physies behind this process remains unclear.
Although the various attempts to simulate disk formation have provided some impressive successes, the essential physics behind this process remains unclear.
Governatoetal.(2002) claim that the low angular momentum of simulated galaxies in previous works owed at least partly to inadequate resolution.
\cite{governato02a} claim that the low angular momentum of simulated galaxies in previous works owed at least partly to inadequate resolution.
In addition. they emphasize the impact that the matter power spectrum can have on galaxy formation through simulations of WDM universes in which the bulge and spheroid components of galaxies are smaller than in CDM models.
In addition, they emphasize the impact that the matter power spectrum can have on galaxy formation through simulations of WDM universes in which the bulge and spheroid components of galaxies are smaller than in CDM models.
These findings are supported by Sommer-Larsenetal.(2002.2003) and Sommer-Larsen&Dolgov(2001).. whose calculations indicate that high resolution. strong stellar feedback. and warm dark matter can produce realistic disk galaxies.
These findings are supported by \citet{slgp02a,slgp03a} and \citet{sld01a}, whose calculations indicate that high resolution, strong stellar feedback, and warm dark matter can produce realistic disk galaxies.
Finally. Abadietal.(2003a) suggest that a crucial ingredient for forming proper disks is an implementation of feedback that can regulate star formation.
Finally, \cite{abadi03a} suggest that a crucial ingredient for forming proper disks is an implementation of feedback that can regulate star formation.
Here. we use a “multiphase model to describe star-forming gas. in Which the ISM consists of cold clouds in. pressure equilibrium with an ambient hot phase (Springel&Hern-quist 2003a).
Here, we use a `multiphase model' to describe star-forming gas, in which the ISM consists of cold clouds in pressure equilibrium with an ambient hot phase \citep{sh03a}.
. Radiative cooling of gas leads to the growth of clouds. which in turn host the material that fuels star formation.
Radiative cooling of gas leads to the growth of clouds, which in turn host the material that fuels star formation.
The supernovae associated with star formation provide feedback by heating the ambient medium and evaporating cold clouds.
The supernovae associated with star formation provide feedback by heating the ambient medium and evaporating cold clouds.
The feedback treatment establishes= a self-regulated cycle for star formation. pressurizing the star-forming gas.
The feedback treatment establishes a self-regulated cycle for star formation, pressurizing the star-forming gas.
Implemented in a subresolution manner. our approach makes it possible to obtain numerically converged results for the star formation rate at moderate resolution.
Implemented in a subresolution manner, our approach makes it possible to obtain numerically converged results for the star formation rate at moderate resolution.
This feature is particularly important for simulations of CDM cosmologies. where early generations of galaxies may be difficult to resolve.
This feature is particularly important for simulations of CDM cosmologies, where early generations of galaxies may be difficult to resolve.
Using this model of star-forming gas. Springel&Hern-quist(2003b) obtained a converged prediction for the history of cosmic star formation that agrees well with observations at redshifts z—4 (Springel&Hernquist20050:&Springel2003)!
Using this model of star-forming gas, \cite{sh03b} obtained a converged prediction for the history of cosmic star formation that agrees well with observations at redshifts $z \le 4$ \citep{sh03b, her03}.
In subsequent work. Nagamineetal.(2003a.b.d) showed that this model also accounts for the observed abundance and star formation rates of damped Lyman-alpha absorbers and Lyman-break galaxies at z3.
In subsequent work, \cite{nsh03a, nsh03b, nshm03} showed that this model also accounts for the observed abundance and star formation rates of damped Lyman-alpha absorbers and Lyman-break galaxies at $z\sim 3$.
In what follows. we study the consequences of our multiphase model for the ISM on the formation of disk galaxies.
In what follows, we study the consequences of our multiphase model for the ISM on the formation of disk galaxies.
First. we employ a high-resolution simulation. to identify realistic disk galaxies in a cosmological context.
First, we employ a high-resolution simulation to identify realistic disk galaxies in a cosmological context.
For one such disk galaxy. we examine its kinematic and£o photometric properties in. detail. demonstrating good agreement with observations of local spirals.
For one such disk galaxy, we examine its kinematic and photometric properties in detail, demonstrating good agreement with observations of local spirals.
Second. we use a set of idealized simulations to study disk formation in individual dark matter halos to isolate the effect of different models for the equation of state of the ISM.
Second, we use a set of idealized simulations to study disk formation in individual dark matter halos to isolate the effect of different models for the equation of state of the ISM.
This analysis demonstrates that the feedback in our multiphase model alters the dynamics by pressurizing the star-forming gas and stabilizing forming disks against fragmentation.
This analysis demonstrates that the feedback in our multiphase model alters the dynamics by pressurizing the star-forming gas and stabilizing forming disks against fragmentation.
We emphasize that this aspect of our modeling does not depend on the details of our prescription for star formation and feedback. but is determined by the effective equation of state for star-forming gas.
We emphasize that this aspect of our modeling does not depend on the details of our prescription for star formation and feedback, but is determined by the effective equation of state for star-forming gas.
Thus. our conclusions should obtain generally. provided that the actual bulk equation of state for the ISM has characteristics similar to those of our description.
Thus, our conclusions should obtain generally, provided that the actual bulk equation of state for the ISM has characteristics similar to those of our description.
In 2.. we present our simulation method.
In \ref{sec:sim:cosmo}, we present our simulation method.
We review our analysis procedure in 3.. and discuss our findings for the structural 4)) and kinematic properties 5)) of one simulated disk galaxy.
We review our analysis procedure in \ref{sec:analysis}, and discuss our findings for the structural \ref{sec:disk:nonkin}) ) and kinematic properties \ref{sec:disk:kin}) ) of one simulated disk galaxy.
In 6.. we deseribe our idealized simulations and their results.
In \ref{sec:sim:halo:intro}, we describe our idealized simulations and their results.
Finally. we conclude and suggest directions for further research in 7..
Finally, we conclude and suggest directions for further research in \ref{sec:conclusions}.
Our simulations were performed using the parallel N-body/smoothed particle hydrodynamics (SPH) code im its "conservative entropy” formulation (SpringelHernquist2002).. to mitigate problems with lack of energy and entropy conservation in older treatments of SPH (e.g. 2003).
Our simulations were performed using the parallel $N$ -body/smoothed particle hydrodynamics (SPH) code in its “conservative entropy” formulation \citep{sh02a}, to mitigate problems with lack of energy and entropy conservation in older treatments of SPH \citep[e.g.][]{her93,oshea03}.
. We adopt a flat ACDM cosmology with cosmological parameters OQ,=0.3. Q420.7. Oy=0.04. and oy=0.9. and set the primordial power spectrum index to 5;=l.
We adopt a flat $\Lambda$ CDM cosmology with cosmological parameters $\Omega_{\rm m}=0.3$, $\Omega_{\Lambda}=0.7$, $\Omega_{\rm b}=0.04$ , and $\sigma_{8}=0.9$, and set the primordial power spectrum index to $n=1$ .
Throughout. we select a value for the Hubble constant of Hy=1007kms?!Mpce! with hiz0.7.
Throughout, we select a value for the Hubble constant of $H_{0} = 100\,h\,{\rm km\, s^{-1} Mpc^{-1}}$ with $h=0.7$.
For our cosmological simulation. we populate a periodic volume of 10 47! Mpe on aside with 144? low-resolution dark matter (LRDM) particles.
For our cosmological simulation, we populate a periodic volume of 10 $h^{-1}$ Mpc on a side with $144^{3}$ low-resolution dark matter (LRDM) particles.
At the center of the box. a 5 // Μρο cubic region is selected as a high resolution region where we replace the LRDM particles with particles of eight-times lower mass.
At the center of the box, a 5 $h^{-1}$ Mpc cubic region is selected as a high resolution region where we replace the LRDM particles with particles of eight-times lower mass.
The initial displacement field is then calculated following a standard "zooming" procedure (Tormen1997;Poweretal. 2003). where small scale perturbations are added appropriately in the high-resolution region.
The initial displacement field is then calculated following a standard “zooming” procedure \citep{tormen97a, power03a}, where small scale perturbations are added appropriately in the high-resolution region.
Note that the high-resolution zone does not target a particular object. with the intent of eliminating bias that could be introduced by selecting halos that may be intrinsically favorable for disk galaxy formation.
Note that the high-resolution zone does not target a particular object, with the intent of eliminating bias that could be introduced by selecting halos that may be intrinsically favorable for disk galaxy formation.
We further split the high-resolution particles into dark matter (HRDM) and gas.
We further split the high-resolution particles into dark matter (HRDM) and gas.
The resulting particle masses of each component are then z/]upwM 1017.. ην=3.02«M10577i. and ma,=4.65«10*47 |.
The resulting particle masses of each component are then $m_{\rm LRDM}=2.79\times 10^{7} h^{-1} M_{\odot}$ , $m_{\rm HRDM}=3.02\times 10^{6} h^{-1} M_{\odot}$, and $m_{\rm gas}=4.65\times 10^{5} h^{-1} M_{\odot}$ .
We set the gravitational softening length for the high-resolutionM particles to 0.65 comoving /r! kpe.
We set the gravitational softening length for the high-resolution particles to 0.65 comoving $h^{-1}$ kpc.
While our 10 77! Mpe box is too small to be fully representative of the >=0 universe. the volume is sufficient for our purposes as our current work does not concern. for example. large-scale correlations of galaxies or the global mass function.
While our 10 $h^{-1}$ Mpc box is too small to be fully representative of the $z=0$ universe, the volume is sufficient for our purposes as our current work does not concern, for example, large-scale correlations of galaxies or the global mass function.
Here. we are interested only in individual. galactic-sized objects which are mostly unaffected by the simulation box size.
Here, we are interested only in individual, galactic-sized objects which are mostly unaffected by the simulation box size.
We include a UV background and radiative cooling and heating in the manner of Katzetal.(1996) and Davéetal. (1999), as well as star formation. supernova feedback. and metal enrichment.
We include a UV background and radiative cooling and heating in the manner of \cite{kwh96} and \cite{dave99}, as well as star formation, supernova feedback, and metal enrichment.
We employ the multiphase model developed by Springel&Hernquist(2003a) to. describe the star-forming gas (seealsoYepesetal.1997:Hultman&Pharasyn 1999).
We employ the multiphase model developed by \cite{sh03a} to describe the star-forming gas \citep[see also][]{yepes97,hp99}.
. Our approach accounts for some of the keyaspects of the multiphase structure of the ISM(McKeewithout spatially resolving the different phases explicitly.
Our approach accounts for some of the keyaspects of the multiphase structure of the ISM\citep{mo77a} without spatially resolving the different phases explicitly.
Instead. a statistical mixture of the phases is computed analytically. taking into account the growth of cold clouds embedded in a supernova-heated ambient phase. the formation of stars out of the cloud material. and the evaporation of clouds in supernova remnants.
Instead, a statistical mixture of the phases is computed analytically, taking into account the growth of cold clouds embedded in a supernova-heated ambient phase, the formation of stars out of the cloud material, and the evaporation of clouds in supernova remnants.
reactions.
reactions.
We reproduce the observed isolopomer ratio CIL/MCCLE by considering the isolopomer-exchange reaction: ΠΟ ΟΕ + IE — CM CLE + I+ 8.1 Ix. if this reaction proceeds with the forward rate coefficient hy > ll!! em? 1.
We reproduce the observed isotopomer ratio $^{13}$ $^{13}$ CCH by considering the isotopomer-exchange reaction: $^{13}$ CCH + H $\rightarrow$ $^{13}$ CH + H + 8.1 K, if this reaction proceeds with the forward rate coefficient $k_f$ $>$ $^{-11}$ $^3$ $^{-1}$.
The isotope ratio is underestimated in our model. which is left for future studies.
The isotope ratio is underestimated in our model, which is left for future studies.
3.
3.
We reproduce the observed. CÉCS/PE CCS. CCS/CP'CS and CCS/P CCS. ratios simultaneously by considering the isotopomer-exchange reaction: MCCS + II CPCS + H+ 17.4 IN. if this reaction proceeds with the forward rate coellicient hy > th em? |
We reproduce the observed $^{13}$ $^{13}$ CCS, $^{13}$ CS and $^{13}$ CCS ratios simultaneously by considering the isotopomer-exchange reaction: $^{13}$ CCS + H $\rightarrow$ $^{13}$ CS + H + 17.4 K, if this reaction proceeds with the forward rate coefficient $k_f$ $>$ $^{-11}$ $^3$ $^{-1}$.
4.
4.
In conclusion. isolopomer fractionation of CCIE and CCS can be due to isolopomer- reactions.
In conclusion, isotopomer fractionation of CCH and CCS can be due to isotopomer-exchange reactions.
Laboratory measurements aud detailed quantum calculations of these isotoponmer exchange reactions are highly Acknowledgements.
Laboratory measurements and detailed quantum calculations of these isotopomer exchange reactions are highly Acknowledgements.
We thank Yoshihiro Osamura for providing the zero-point energies of CE'CS and PCCS. and valuable discussions.
We thank Yoshihiro Osamura for providing the zero-point energies of $^{13}$ CS and $^{13}$ CCS, and valuable discussions.
We also thank Riccardo Tarroni for providing the zero-point energies of CMCLE and. MCCLL
We also thank Riccardo Tarroni for providing the zero-point energies of $^{13}$ CH and $^{13}$ CCH.
We would like to thank the anonymous releree for the helpful comments to improve (he manuscript.
We would like to thank the anonymous referee for the helpful comments to improve the manuscript.
This work is supported by a erant-in-aid [or scientific research (21244021) and Global COE program |Foundation of International Center for Planetary Science" (G11) of the Ministry. of Education. Culture. Sports. Science aud Technology of Japan (MEXT).
This work is supported by a grant-in-aid for scientific research (21244021) and Global COE program "Foundation of International Center for Planetary Science" (G11) of the Ministry of Education, Culture, Sports, Science and Technology of Japan (MEXT).
OGLE lighteurve reveals nothing of interest and it is not a detectable LR source in the 2ALASS data.
OGLE lightcurve reveals nothing of interest and it is not a detectable IR source in the 2MASS data.
Hence it cannot be a strong contender for the counterpart to AN τοῦ.
Hence it cannot be a strong contender for the counterpart to AX $-$ 733.
On the other hand. Object 512 has V=154 and a significant LR flux at /=15.3.
On the other hand, Object 512 has $V=15.4$ and a significant IR flux at $J=15.3$.
Both of these make it look like a classic counterpart to a οX-ray binary system.
Both of these make it look like a classic counterpart to a Be/X-ray binary system.
νο compare this object to another SMC X-ray. pulsar svsten. IWGA 7226 (Buckley ct al. 2001). we find it is extremely similar.
If we compare this object to another SMC X-ray pulsar system, 1WGA $-$ 7226 (Buckley et al, 2001), we find it is extremely similar.
In. IWOA 7226 we have compared to 0.03 in Object 512. and (JA)=0.62 compared to 0.51 in Object 512.
In 1WGA $-$ 7226 we have $(B-V)=-0.06$ compared to $-0.03$ in Object 512, and $(J-K)=0.62$ compared to 0.51 in Object 512.
The οY) value found for many other SAIC counterparts to De/X-ray systems is ~0.25 (a combination of extinction to the SAIC plus local extinction due to cirumstellar material).
The $E(B-V)$ value found for many other SMC counterparts to Be/X-ray systems is $\sim0.25$ (a combination of extinction to the SMC plus local extinction due to cirumstellar material).
Applying this to the observed. values for Object 512 given in Table l leads to an identification for the spectral type of BOLLL-V.
Applying this to the observed values for Object 512 given in Table 1 leads to an identification for the spectral type of B0III-V.
Thus even before one considers the source. one is lecl inexorably to Object 512 being the prime candidate for the optical counterpart to AN 0051-733.
Thus even before one considers the source, one is led inexorably to Object 512 being the prime candidate for the optical counterpart to AX J0051-733.
The presence of a convincing source at the same position adds significant extra weight to this conclusion.
The presence of a convincing source at the same position adds significant extra weight to this conclusion.
The optical spectrum of Object 512. presented. in ligure 2. is no later ancl perhaps slightly. earlier than the comparison standard.
The optical spectrum of Object 512 presented in Figure \ref{os} is no later and perhaps slightly earlier than the comparison standard.
From the colours presented in Table 1 and assuming (2Vy=0.26 (Wegner 1904). this results in an extinction value of £(2VV)=0.23. which confirms the number used above in interpreting just the photometry.
From the colours presented in Table 1 and assuming $(B-V)_{0} = -0.26$ (Wegner 1994), this results in an extinction value of $E(B-V)=0.23$, which confirms the number used above in interpreting just the photometry.
Assuming standard reddening. 24)=0.71 and therefore. assuming a distance modulus to the SAIC (ALοἱ)=15.0. the absolute magnitude for Object 512 is Ady=4.2. which is in rather good agreement with a spectral type in the Bo-D30.5V range.
Assuming standard reddening, $A_{V} = 0.71$ and therefore, assuming a distance modulus to the SMC $(M-m)=18.9$, the absolute magnitude for Object 512 is $M_{V} = -4.2$, which is in rather good agreement with a spectral type in the B0-B0.5V range.
The strong sinusoidal optical modulation of Object 512 is challenging to interpret in terms of a traditional Be/X-ray binary model.
The strong sinusoidal optical modulation of Object 512 is challenging to interpret in terms of a traditional Be/X-ray binary model.
Firstly. the expected. binary period of AX 738 based on the Corbet diagram (Corbet. 1986) is 100200d.
Firstly, the expected binary period of AX $-$ 733 based on the Corbet diagram (Corbet, 1986) is 100–200d.
Secondly. a binary period. of just. 14d. involving a Be star implies an extremely tight orbit the Ixeplerian orbital radius would be ~14 solar radii and the BO star has a size of S solar radii.
Secondly, a binary period of just 1.4d involving a Be star implies an extremely tight orbit – the Keplerian orbital radius would be $\sim14$ solar radii and the B0 star has a size of $\sim8$ solar radii.
Phirdly. if the period is really decreasing at a rate of 13.5 s/vear then this implies (LIuang 1963) a mass transfer of 10.7AZ. /vear for mass transfer between an 18M. Be star and a 1.42. neutron star which is not only much larger than that typically observed in LIAIND systems (αςLO7AL. fvear in most cases). but would also imply a much higher X-ray luminosity unless the accretion mechanism is extremely incllicient at. converting eravitational potential into X-rays.
Thirdly, if the period is really decreasing at a rate of 13.5 s/year then this implies (Huang 1963) a mass transfer of $10^{-5}\: M_{\odot}$ /year for mass transfer between an $M_{\odot}$ Be star and a $M_{\odot}$ neutron star – which is not only much larger than that typically observed in HMXB systems $\la 10^{-8}\: M_{\odot}$ /year in most cases), but would also imply a much higher X-ray luminosity unless the accretion mechanism is extremely inefficient at converting gravitational potential into X-rays.
Alass transfer rates of this magnitude are deduced to exist in the EB binary svstem 2 Lyrae which is changing its 13d orbital period at a rate of 198/vear (van Llanime. Wilson CGuinan 1995).
Mass transfer rates of this magnitude are deduced to exist in the EB binary system $\beta$ Lyrae which is changing its $\sim13$ d orbital period at a rate of 19s/year (van Hamme, Wilson Guinan 1995).
In this case the change is to a longer period with the mass transferring from the smaller BG-S star to the more massive Be star.
In this case the change is to a longer period with the mass transferring from the smaller B6-8 star to the more massive Be star.
In our case. the mass would be [owing in the opposite direction. i.e. from the more massive object to à less massive one.
In our case, the mass would be flowing in the opposite direction, i.e. from the more massive object to a less massive one.
The optical lightcurve of 3 Lyrae is similar to the one presented here for Object 512. but with the notable cillerence that in; Lyrae the two minima are not of the same depth.
The optical lightcurve of $\beta$ Lyrae is similar to the one presented here for Object 512, but with the notable difference that in $\beta$ Lyrae the two minima are not of the same depth.
1n fact the symmetry of the light curve is much more sugeestive of a \W UMa type system.
In fact the symmetry of the light curve is much more suggestive of a W UMa type system.
Unfortunately. the observed. period. of 1-4d. is much greater than any. such reported system in the SAIC (Rucinski 1997).
Unfortunately, the observed period of 1.4d is much greater than any such reported system in the SMC (Rucinski 1997).
Phe maximum observed. period is O.Scl and our period is well off the end of he distribution.
The maximum observed period is 0.8d and our period is well off the end of the distribution.
In addition. it is perhaps worth noting tha he predicted (Vo£) colour obtained from the distribution of such objects and our binary period of 14d is. |0.026. bu rom ‘Table Lit can be seen that the observed (V£) for Object 512 is 0.17.
In addition, it is perhaps worth noting that the predicted $(V-I)$ colour obtained from the distribution of such objects and our binary period of 1.4d is $+0.026$, but from Table 1 it can be seen that the observed $(V-I)$ for Object 512 is 0.17.
Even allowing for interstellar extinction his further aces to it being unlikely that this svstem is of his class.
Even allowing for interstellar extinction this further adds to it being unlikely that this system is of this class.
The possibility of a blended: variable star plus Be star can be considered.
The possibility of a blended variable star plus Be star can be considered.
For example. a chance superposition of De star (to give the observed. colours) plus Cepheicl or Rik Lyrae (to give the optical modulation).
For example, a chance superposition of Be star (to give the observed colours) plus Cepheid or RR Lyrae (to give the optical modulation).
Llowever. all of these models can be ruled out because of either the magnitude of he period. or the depth of modulation. or the shape of the ightcurve.
However, all of these models can be ruled out because of either the magnitude of the period, or the depth of modulation, or the shape of the lightcurve.
Interestingly the optical modulation. is) somewha similar to the short. periodic modulation seen. by Dalona 992) in Be stars in the cluster NCC30 in the SAIC.
Interestingly the optical modulation is somewhat similar to the short periodic modulation seen by Balona (1992) in Be stars in the cluster NGC330 in the SMC.
In this case Dalona attributes this modulation to surface features on the rapidly rotating objects.
In this case Balona attributes this modulation to surface features on the rapidly rotating objects.
However. how the period of such objects could change on a timescale of vears is not clear. unless the star is in a very wide binary system.
However, how the period of such objects could change on a timescale of years is not clear, unless the star is in a very wide binary system.
Lt is possible that the data in Figure 6. could be fitted to ~10 Eear sinusoidal modulation. but then the orbit of the neutron star would be so distant from the Be star that it hard to see how accretion could ever occur.
It is possible that the data in Figure \ref{per} could be fitted to $\sim$ 10 year sinusoidal modulation, but then the orbit of the neutron star would be so distant from the Be star that it hard to see how accretion could ever occur.
In addition X-ray outbursts have been detected 3 times over 2 vears from this svstem (Lavcock. private communication) making such a long orbit unlikely.
In addition X-ray outbursts have been detected 3 times over 2 years from this system (Laycock, private communication) making such a long orbit unlikely.
Perhaps further optical data may clarify exactly what the shape of the period change is on such timescales.
Perhaps further optical data may clarify exactly what the shape of the period change is on such timescales.
We are loft with no convincing traditional scenario to explain all the observational data.
We are left with no convincing traditional scenario to explain all the observational data.
It is very hard to see how the orbital period change seen in Figure ο could. possibly be caused by mass loss from a normal DO star at a rate of 10.7 Al. fvear.
It is very hard to see how the orbital period change seen in Figure \ref{per} could possibly be caused by mass loss from a normal B0 star at a rate of $10^{-5}$ $M_{\odot}$ /year.
One other possibility perhaps worth considering is that AX 733 is a triple system 3e star plus another star in a tight. 14d. orbit. and the neutron star in a highly eccentric 2000. orbit around the pair.
One other possibility perhaps worth considering is that AX $-$ 733 is a triple system – Be star plus another star in a tight 1.4d orbit, and the neutron star in a highly eccentric $-$ 200d orbit around the pair.
Such a system could not only be intrinsically very stable since most of the mass is concentrated in the inner binary pair. but the transfer of angular momentum from the inner binary to the orbit of the neutron star might also explain the evolution of the orbital period.
Such a system could not only be intrinsically very stable since most of the mass is concentrated in the inner binary pair, but the transfer of angular momentum from the inner binary to the orbit of the neutron star might also explain the evolution of the orbital period.
Eeeleton Ixiseleva (1995) derive a critical parameter AUUU for a stable triple svstem. which is the period. ratio
Eggleton Kiseleva (1995) derive a critical parameter $X_{o}^{min}$ for a stable triple system, which is the period ratio
objects.
objects.
It is known that one property of the accretion process that characterizes PMS stars is temporal variability, which can be traced to changes both in the continuum and more importantly in the emission lines (e.g.Herbst1986;Hartiganetal.1991;Nguyenetal. 2009).
It is known that one property of the accretion process that characterizes PMS stars is temporal variability, which can be traced to changes both in the continuum and more importantly in the emission lines \citep[e.g.][]{herb86,hart91,ngu09}.
. As mentioned in Section 3.2,, SB04 published list of sources with Ha emission in 33603 derived froma SSO and WFPC2 observations.
As mentioned in Section \ref{sec_lett}, SB04 published a list of sources with $\alpha$ emission in 3603 derived from SSO and WFPC2 observations.
We searched for the counterparts of these objects in our observations by cross-correlating our photometric catalogue with that of SB04.
We searched for the counterparts of these objects in our observations by cross-correlating our photometric catalogue with that of SB04.
The stars detected by SB04 on the SSO data are bright and are all saturated in our images.
The stars detected by SB04 on the SSO data are bright and are all saturated in our images.
On the other hand, the WFPC2 catalogue of SB04 contains 96 sources selected on the basis of their Ho index, of which 67 have a counterpart in our catalogue.
On the other hand, the WFPC2 catalogue of SB04 contains 96 sources selected on the basis of their $\alpha$ index, of which 67 have a counterpart in our catalogue.
As for the 29 missing objects, ten fall in the inner 10" of the cluster center, where saturation in our images makes star detection very difficult.
As for the 29 missing objects, ten fall in the inner $10\arcsec$ of the cluster center, where saturation in our images makes star detection very difficult.