D. Cordier^{1,2}  Y. Lebreton^{1}  M.J. Goupil^{1}  T. Lejeune^{4}  J.P. Beaulieu^{5}  F. Arenou^{1}
1  DASGAL, CNRS UMR 8632,
Observatoire de ParisMeudon, DASGAL,
92195 Meudon Principal Cedex, France
2  École Nationale Supérieure de Chimie de Rennes,
Campus de Beaulieu, 35700 Rennes, France
3  Observatorio Astronomico, Universidade de Coimbra,
Santa Clara 3040 Coimbra, Portugal
4  I.A.P., 98bis boulevard Arago, 75014 Paris, France
Received 5 March 2002 / Accepted 6 June 2002
Abstract
The main purpose of this paper is to investigate the possible existence of a
metallicity dependence of the overshooting from main sequence star turbulent cores.
We focus on objects with masses in the range

.
Evolutionary time scale ratios are compared with star number
ratios on the main sequence. Star populations are synthesized using grids of evolutionary
tracks computed with various overshooting amounts. Observational material is provided by
the large and homogeneous photometric database of the OGLE 2 project for the Magellanic clouds.
Attention is paid to the study of uncertainties: distance modulus, intergalactic and interstellar
reddening, IMF slope and average binarity rate. Rotation and the chemical composition
gradient are also considered. The result for the overshooting distance is
(
Z_{0}=0.004) and
(
Z_{0}=0.008) suggesting a possible dependence of the extent of
the mixed central regions with metallicity within the considered mass range. Unfortunately it is
not yet possible to fully disentangle the effects of mass and chemical composition.
Key words: convection  stars: evolution, interiors
Extensive convective phenomena occur in the cores of main sequence stars with masses above about 1.2 (for galactic chemical composition). In standard models, convection is crudely modeled with the wellknown Mixing Length Theory of BöhmVitense (1958) (hereafter MLT) and the core extension is determined according to the Schwarzschild criterion. The Schwarzschild limit is the value of the radius where the buoyancy force vanishes. However, inertia of the convective elements leads to an extra mixing above the Schwarzschild limit, called "overshooting'' and is usually expressed as a fraction of the pressure scale height. Several theoretical works (for a review see Zahn 1991) give arguments in favor of such additional mixing. Many laboratory experiments show evidence for overshooting (see Massaguer 1990). Although overshooting can occur below an external convective zone (see Alongi et al. 1991), this paper is exclusively concerned with core overshooting.
One of the first empirical determinations of convective core overshooting was obtained by Maeder & Mermilliod (1981) who used a set of 34 galactic open clusters and fitted the main sequence width with an additional mixing of about 2040% in mass fraction. Mermilliod & Maeder (1986) derived an overshooting amount of about for solarlike chemical composition and for a 915 range. Stothers & Chin (1991) derived an overshooting amount < for Pop. I stars using the metalenriched opacity tables published in Rogers & Iglesias (1992).
During the last decade, many evolutionary model grids have been computed with an
overshooting amount equal or close to
:
e.g.
Charbonnel et al. (1996) or Bertelli et al. (1994). This second team (Padova group) uses
a formalism (see Bressan et al. 1981) slightly different from the Geneva
team one (e.g. see Schaller et al. 1992). Generally the same overshooting amount
is used whatever the metallicity and mass are.
Figure 1: a) CMDiagram for the OGLE 2 SC 1 field, 10% of the data have been plotted for sake of clarity. b) Standard deviation of the measurements: versus magnitude V, the dotdashed line indicates the limitvalue considered (0.015). c) The resulting CMdiagram after selection (all the fields within SMC have been plotted).  
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KozhurinaPlatais et al. (1997) obtained for the galactic cluster NGC 3680 (solar metallicity) with the isochrone technique. This method consists of fitting the cluster CMD features (particularly the turnoff position) with model isochrones. Iwamoto & Saio (1999) compared evolutionary models with observations of three binary systems: V2291 Oph, Aur and And ("binary system'' technique). The authors adjusted either the helium content or the overshooting parameter to get a better fit to observations. The best results were obtained with a moderate overshooting amount ( ). For supersolar metallicity ( Z_{0}= 0.024) Lebreton et al. (2001) derived from the modeling of the Hyades cluster turnoff.
Maeder & Mermilliod (1981) have suggested an overshooting increasing with mass within the studied range of 26 which is also found by Schröder et al. (1997) with a study of binary systems. According to their results, the overshooting should increase from for 2.5 to for 6.5 . With a similar study Ribas et al. (2000) also found a mass dependence.
The question of a metallicity dependence must also be addressed. Ribas et al. (2000)'s results suggest a slight metallicity dependence for a stellar mass around 2.40 (see their Table 1). The more metal poor star SZ Cen (in mass fraction: Z_{0}=0.007) is satisfactorily modelled with an overshooting distance and objects with Z_{0} ranging between 0.015 and 0.020 seem to have an overshooting around 0.2 . Keller et al. (2001) have recently explored the dependence of overshooting with metallicity by means of the isochrone technique using isochrone grids from the Padova group. Their study involves HST observations of four clusters: NGC 330 (SMC), 1818, 2004 and 2001 (LMC). Keller et al. (2001) find the best fit (with respect to age and overshooting) for an overshooting amount which is equivalent to in the Geneva formalism ( ). In this paper, we carry out an independent study of a possible metallicity dependence of overshooting with a technique which differs from the "binary system'' (Ribas et al. 2000; Andersen 1991) and "isochrone'' techniques. Our method is based on starcount ratios, with comparisons between observational material and synthetic population results in colormagnitude (CMD) diagrams. We are then led to discuss several points: particularly distance modulus, reddening and binarity rate. If the dependence of overshooting on metallicity (or mass) was thereby to be firmly assessed, it would then be a challenge to understand its physical origin.
We are concerned with a metallicity range relevant to the Magellanic Clouds and take advantage of the homogeneous OGLE 2 data, which provide color magnitude diagrams for stars in the Small Magellanic Cloud (hereafter SMC) and in the Large Magellanic Cloud (hereafter LMC). On the theoretical side, we estimate the number of stars from evolutionary model sequences computed with different amounts of overshooting. From these data sets and using evolutionary models with intermediate and low metallicity, we estimate the overshooting value during the main sequence in the SMC and LMC for a stellar mass in the range 2.5  .
In Sect. 2 we describe the observational data involved in this work. Section 3 is devoted to the method used: data selection and star counting. Section 4 gives the main features of our population synthesis procedure. Section 5 is devoted to astrophysical inputs, and Sect. 6 to results and effects of uncertainties. Section 7 discusses the results. It must be emphasized that we determine in fact the extent of the inner mixed core region which can be due either to true overshooting or to another process such as rotation; some observational evidence exists about correlation between metallicity and (see Venn et al. 1999). The problem of rotation is briefly discussed in Sect. 7. Finally, Sect. 8 gives some comments and concluding remarks. An appendix has been added to provide details about the population synthesis algorithm and error simulations.
The observational data set considered here has been obtained by the Optical Gravitational Lensing Experiment (OGLE hereafter) consortium during its second operating phase (for more details and references the reader can consult URL: http://www.astrouw.edu.pl/ogle/).
We have downloaded the SMC data described in Udalski et al. (1998). The data used in this paper are from the postApr. 8, 2000 revision. The SMC is divided into 11 fields (labeled SC1 to SC11) covering ; each field contains between 100 000 and 350 000 objects. For each object several quantities are available: equatorial coordinates, BVI photometry and associated standard errors , and . This database has the great advantage of being extensive and very homogeneous.
The LMC data are described in Udalski et al. (2000). The BVI map of the LMC is composed of 26 fields (SC1 to SC26) in the central bar of the LMC. The dataset includes photometry and astrometry for about 7 million stars over a 5.7 square degree field.
As shown in Fig. 1b, the standard error on Vmagnitude, , increases with the magnitude. This is also true for B or Imagnitudes. Hence the errors on (BV) or (VI) colors rapidly increase and reach values as large as 0.2 mag around a Vmag 20: this is of the same order as the Main Sequence width.
As we are interested in the MS structure and as we must minimize error effects while keeping quite good statistics, we have chosen to take into account only data with and (or ) lower or equal to 0.015 mag, leading to a maximum error on color of 0.02 mag. The value of 0.015 mag appears to be an optimal choice maintaining a good statistics with photometric errors remaining small compared with the MS width. Figure 1 sketches the proposed selection process and displays differences between the entire ColorMagnitude Diagram (Fig. 1a) and the final diagram (Fig. 1c): obviously, the remaining data are those corresponding to lower magnitudes.
This selection process leaves 4700 objects on the SMC MS (over a total of more than 2.2 millions objects) in the BV system (1100 objects in the VI system) and 4000 objects on the LMC MS (over a total of more than 7.2 millions objects) in the BV system (1600 objects in the VI system). As we can see, the BV system presents more favorable statistics, therefore in the following we will work only with this set of bands.
Tables 4 from Udalski et al. (1998, Udalski et al. 2000) indicate that completeness for
should be better than about 99% for the SMC; and
should be around
depending on the field crowding for the LMC.
Figure 2: a) Data from SMC with mag and box definitions ( N_{1}+N_{2}= 4653 and N_{2}/N_{1}= 1.08), b) the same for the LMC case ( N_{1}+N_{2}= 4113 and N_{2}/N_{1}= 1.01).  
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As the absolute number of stars arriving on the ZAMS per unit of time for a given mass is unknown, we rather compute star count ratios. To count stars, we first define an area in the CMdiagram. As we are interested in the MS structure, we choose a region which contains the main sequence "bulge'' revealed after the data selection process (see Fig. 2a) with the most convenient geometrical shape: a "parallelogram'' (for automatic count purpose). A couple of opposite sides (AB and DE in Fig. 2a) are chosen to be more or less "parallel'' to the main sequence axis.
In the CMDiagram, main sequence stars evolve from the blue to the red side. The MS width is mainly an evolutionary effect connected to a characteristic time scale (time spent by a star on the Main Sequence). The distribution of the objects within the Main Sequence should be related to this time scale. Therefore we divide our parallelogram into two regions called "box 1'' and "box 2'' (see Fig. 2a) where the respective numbers of objects N_{1} and N_{2} are similar ( ). This ratio is taken as an observational constraint and it will enable us to discriminate between theoretical grids of evolutionary tracks computed with various overshooting amounts.
We now turn to the method used to build a synthetic stellar sample comparable to the OGLE 2 ones (after selection) from evolution simulation outputs.
Our evolutionary models are built with the 1D Henyey type code CESAM^{} (see Morel 1997) in which we brought several improvements. Applying modern techniques like the projection of the solutions on Bspline basis and automatic mesh refinements, CESAM allows robust, stable and highly accurate calculations. We use as physical inputs:
In order to compare theoretical results to observational data, conversions are needed. Transformations of the theoretical quantities, ( , ) into absolute magnitudes and colors are derived from the most recent version of the Basel Stellar Library (BaSeL, version 2.2), available electronically at ftp://tangerine.astro.mat.uc.pt/pub/BaSeL/. This library provides colorcalibrated theoretical flux distributions for a large range of fundamental stellar parameters, (2000 to 50 000 K), (1.0 to 5.5 dex), and (5.0 to +1.0 dex). The BaSeL flux distributions are calibrated on the stellar UBVRIJHKL colors, using:
In contrast with "classical'' works on population synthesis where the
CMD as a whole is simulated, we construct a small part
of the CMD: the area containing the brighter MS stars. In this way the task is simplified.
Artificial stellar samples have been generated from our evolutionary tracks with a specially
designed population synthesis code CReSyPS
^{}.
In our framework the main hypothesis is that the Star Formation Rate (SFR) is constant during the time scales involved here: i.e. a few hundred megayears. So for a given mass the number of observed stars (i.e. those corresponding to a given evolutionary track) must be proportional to the time scale of the main sequence. We assume that the SFR is constant in time and mass (equal for all masses in the range explored in this work), if we note r the SFR: represents the mean time elapsed between two consecutive star births. For the observational star samples, is unknown but the objects numbers are available. We choose to get similar total star numbers in boxes 1 and 2 (i.e. N_{1}+N_{2}) both in the synthetic CMD and observational diagram. We point out that the ratios N_{2}/N_{1} are not sensitive to the value chosen.
The evolutionary track grids scan a mass range between 2.5
and 25
from the ZAMS to
covering the entire
box ranges in color and magnitude (defined in Sect. 3.2).
The mass step is increasing from 0.5
around 3
stars to
5
above 15
.
Several overshooting amounts have been used from 0.0
to 0.8. CReSyPS
treats the photometric errors by simulating OGLE 2 ones
(see Appendix A) which is very important for our purposes.
Our algorithm requires the knowledge of some input parameters:
distance modulus, reddening
and absorption, binarity rate, Initial Mass Function (hereafter IMF) slope and photometric errors.
We summarize here the main steps of the algorithm:
Our code intensively uses a random number generator. We have chosen an algorithm insuring a very large period about (program "ran2'' from Press et al. 1992), which is much larger than the number of synthetized objects.
As a result, examples of synthetic samples generated by CReSyPS
are displayed
in Fig. 3 where the influence of overshooting is shown for both clouds.
Figure 3: Synthetic CMDiagrams for SMC and LMC chemical compositions, panels a) and b) are for the SMC with two overshooting amounts: and , panels c) and d) are for the LMC with: and respectively. For panel a) only 30% of the synthetic objects have been displayed for clarity purpose and corresponding evolutionary tracks have been plotted. For other panels: b) d) the number of displayed objects has not be reduced. In all cases the total number of stars in the boxes used for calculations is close to the observational one: a) N_{1}+N_{2}= 4403 ( N_{2}/N_{1}= 0.65), b) N_{1}+N_{2}= 5011 ( N_{2}/N_{1}= 1.59), c) N_{1}+N_{2}= 3582 ( N_{2}/N_{1}= 0.94), d) N_{1}+N_{2}= 4717 ( N_{2}/N_{1}= 1.13). We recall that empirically we got for the SMC N_{1}+N_{2}= 4653 ( N_{2}/N_{1}= 1.08) and for the LMC N_{1}+N_{2}= 4113 ( N_{2}/N_{1}= 1.01). The ratios N_{2}/N_{1} obtained theoretically are rather independant from the value of N_{1}+N_{2} in the synthetic CMDiagram, for instance in the case of the panel b) we got N_{2}/N_{1}= 1.60 with N_{1}+N_{2}= 6399. The cloud of dots in panel b) ( , SMC) calls a comment: it appears to be bimodal, i.e. showing overpopulated regions around mag and mag. Indeed for high overshooting values the main sequence of masses as small as can reach generating with their large evolutionary time scale an overpopulated region. Moreover the binarity shifts a part of this population to a 0.8 mag brighter region, with a binarity rate , this "bimodal effect'' disappears.  
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We have chosen to model the SMC depth with a gaussian distribution of
distances around
with a standard deviation:
We have to distinguish: foreground reddening (due to material in Milky Way) and internal reddening with an origin into the Cloud itself. These quantities are expected to change along the lineofsight. Here we model the total reddening as taking into account its nonuniformity. From the literature, we derive estimations for the mean value and the dispersion of , objecttoobject variations can then be simulated.
We now discuss reddening determinations for the SMC and LMC. From a study of spectral properties of galactic nuclei behind the Magellanic Clouds, Dutra et al. (2001) have evaluated the foreground and background reddenings for both Clouds. For the LMC, they found an average spectroscopic reddening of mag. The uncertainties essentially come from the determination of the stellar populations belonging to background galaxies: in the case of LMC, when Dutra et al. (2001) consider only red population galaxies, they find mag, which gives an idea of the global uncertainty on E(BV), which should be around 0.02  0.03 mag (about ). For the SMC Dutra et al. (2001) find mag. The OGLE 2 project provides reddening for each Cepheid star discovered in both Clouds. OGLE values are: (SMC) and (LMC). In Fig. 4 we have displayed the histogram of values from Dutra et al. (2001) and OGLE group. OGLE data have a better statistics with respectively 1333 (SMC) and 2049 (LMC) objects, against 14 (SMC) and 22 (LMC) for Dutra et al. (2001). Dutra et al's data are systematically less red; this could be inherent to their method: they observed objects behind Clouds and observations are easier through the more transparent regions of the clouds.
In addition, Oestreicher et al. (1995) have
determined the reddening for 1503 LMC foreground
stars with a UBV photometry based method:
mag,
a quite low value because it is related to foreground stars.
It shows a spread (0.02) similar to the OGLE 2 one.
Oestreicher et al.'s (1995)
distribution is in very good agreement (see Fig. 4b) with Dutra
et al.'s one, which tends to confirm that Dutra et al.'s result could be
underestimated (Dutra et al.'s results are supposed to take account foreground and
internal reddeing). Therefore in the case of LMC, we prefer to retain the OGLE average value
for purpose of consistency:
Figure 4: Relative number of stars versus reddening from Dutra et al. (2001) (dashed curve), OGLE 2 experiment Cepheid catalogue: see Udalski et al. (1999) and Udalski et al. (1999), and data read in Fig. 24 of Oestreicher et al. (1995) (dotdashed curve). a) Data for SMC, b) data for LMC.  
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Evaluating the average binary rate in objects as extended as the Magellanic Clouds is not easy. Locally (i.e. within a particular area of the galaxy) this multiplicity rate depends  at least  on two factors: (1) the star density and the kinematics of the objects which influence the encounter probability; (2) the initial binary rate (relative number of binaries on the ZAMS). Within the Magellanic Clouds, the binary rate likely varies over a wide range and we only consider its spatial average value .
Ghez (1995) finds in the solar neighbourhood that for main sequence stars and young stars the binary rate ranges between 0.10 and 0.50 (it peaks at ). Therefore we tested the effects of binarity for these two extreme values.
In our population synthesis code, binaries are taken into account with a uniform probability for the mass ratio q=M_{2}/M_{1} (in the considered mass range).
The IMF has been extensively discussed by many authors. Toward both Galactic poles
and within a distance of 5.2 pc from the Sun, Kroupa et al. (1993) found a mass function:
with
for stars more
massive than 1 .
In the LMC, Holtzman et al. (1997) inferred  from HST observations  a value
consistent with the Salpeter (1955) one:
.
At very
low metallicity, Grillmair et al. (1998) observed the Draco Dwarf spheroidal Galaxy (
)
with the HST. They concluded that the Salpeter IMF slope remains valid in the Magellanic Clouds and
we have chosen:
For a given mass, the Star Formation Rate (SFR) represents the number of stars "created'' per unit of time. Vallenari et al. (1996) have studied three stellar fields of the LMC and have found a time scale of about 24 Gyr for the "bulk of star formation''. We therefore make the reasonable assumption that the SFR remained quite constant during the short galactic period relevant for this work, i.e. for the last 300 Myr. The SFR involved here is an average value over each cloud.
As a first step we choose the mean values for each astrophysical input
(discussed in Sect. 5), this yields for the LMC the following overshooting:
Figure 5: Overshooting determinations for SMC (panels a d)) and LMC (panels e h)). The influence of distance modulus, reddening and IMF slope are considered for each cloud: continuous lines correspond to central values of these parameters discussed in Sect. 5 and dashed lines to the associated error bars. Inferred values of (and uncertainties) are given on each figure. For c) and g) panels, dotted line is for and dotdashed for (see text).  
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For the SMC, using the mean value of each astrophysical inputs we obtain
(see also Fig. 5):
Roxburgh's criterion (Roxburgh 1989) is a very general constraint on the size
of the convective core. It is written as an integral formulation over the stellar
core radius:
The viscous dissipation is unknown but the integral constraint is satisfied for larger value when . Hence, neglecting the dissipation by setting provides the maximum possible extent of the convective core which can be considered as the upper limit for overshooting. Evolutionary tracks have been calculated, using Roxburgh's criterion, for a representative mass of 6 and SMC and LMC metallicities. The equivalent overshooting amount (EOA), given in Table 1, is the time weighted average overshooting distance along the evolutionary tracks, expressed in pressure scale height.
In both cases (LMC and SMC), Roxburgh's criterion predicts a maximum value (i.e. neglecting
viscous dissipation) around
(see Table 1) independent from Z_{0}.
Our determinations  i.e.
and
 therefore are compatible with the theoretical upper limit given by the
Roxburgh's criterion.
Metallicity Z_{0}  0.004 (SMC)  0.008 (LMC) 
Average EOA  0.6  0.6 
In addition to convection, rotation is an other important phenomenon inducing mixing through shear effects and other instabilities. For instance Venn (1999) finds surface abundance variations in SMC A supergiants that could be explained by some kind of mixing related to rotation.
Taking account of the rotational effect brings new important unknown features: (1) the value distribution and (2) the distribution for the considered stellar population. Both features remain unconstrained by observational studies.
In addition, stellar rotation involves many effects and physical processes that are nontrivial to include in modern evolutionary codes. Talon et al. (1997) show that (see their Fig. 5) a rotating 1Dmodel with an initial surface velocity of 300 km s^{1} leads to a main sequence track equivalent to an overshooting model using . Despite great theoretical efforts, a free parameter remains for horizontal diffusivity in Talon et al. (1997) treatment of rotational mixing (see Zahn 1992).
Rotation changes the global shape of an evolutionary track, through two distinct effects: (1) the material mixing inside the inner part of the star which brings more fuel into the nuclear burning zones like overshooting, (2) the effective surface gravity modification leading to color and magnitude changes (which depend on the angle between the lineofsight and the rotational axis). In their Fig. 6, Maeder & Meynet (2001) show the influence of rotation on evolutionary tracks for low metallicity objects ( Z_{0}=0.004). These tracks have been calculated taking into account: (1) an "average effect'' on surface, (2) the internal mixing. These tracks are very similar to those calculated with different overshooting amounts values.
An additional effect which needs to be discussed here is
the surface effect: modifications of colors and magnitudes of MS stars due to rotation
(in absence of any mixing phenomenon) have been studied by Maeder & Peytremann (1970) with uniformly rotating models.
Their Table 2 gives expected changes of M_{V} and (BV) as a function of
(angular
velocity expressed in breakup velocity unit) and
(this latter ranges from 0 to 457 km s^{1},
for a 5
star). In this table standard deviation for M_{V} and (BV),
are:
mag and
mag. Therefore the rotation effect has roughly
the same order of magnitude than present uncertainties on magnitudes and colors.
Figure 6: Relative number of stars ( versus for Wolff & Simon (1997) objects with magnitudes in our studied range.  
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We have sofar assumed a uniform chemical composition. The chemical composition may vary inside each Magellanic Cloud. The existence of an abundance gradient in the Clouds is still debated and spectroscopic measurements with a statistics as large as the statistics of OGLE 2 data are not available. In their Table 4, Luck et al. (1998) give spectroscopic determinations of for 7 SMC Cepheids and 10 LMC Cepheids. For SMC data, the standard deviation is dex leading to negligible variations for the heavy elements mass fraction Z_{0}. Therefore the SMC can be considered as chemically homogeneous for our purpose. For LMC, Luck et al. (1998) find a standard deviation dex giving . From evolutionary tracks of typical mass (6 ) and an overshooting of , changing Z_{0} from 0.007 to 0.01 has a negligible effect on magnitude and an effect of 0.003 mag on color, which is largely lower than the photometric errors. We conclude that  in the light of the present knowledge  the chemical composition gradient does not change our results significantly.
From the investigation of young clusters in the Magellanic Clouds, Keller et al. (2001) did not find any noticeable overshooting dependence with metallicity. They obtained for NGC 330 ( ) , which is compatible with our determination for the SMC: . For NGC 2004 ( ) Keller et al. (2001) got ; while for similar metallicity we derived which is also compatible with Keller et al.'s result. One can note that masses involved in our simulations (average mass of with a standard deviation of ) are higher than the Keller et al. (2001) one (terminus masses in the range for the four clusters). Keller et al. (2001) do not discuss the influence of the uncertainty on distance modulus of the clusters and use mag and mag.
Ribas et al. (2000) derive overshooting amounts from evolutionary models of galactic binary systems.
For SZ Cen (
)
they find
which is close to our value for the LMC, but the mass of SZ Cen is
and
some mass effect
cannot be avoided, therefore any comparison with the present results must be considered with care.
Figure 7: Overshooting parameter versus metallicity Z_{0} from various sources. Open triangles represent results from Ribas et al. (2000) for SZ Cen ( ) error bars have been indicated, arrows mean that the derived value is a minimum. The open square shows a result from KozhurinaPlatais et al. (1997) for the galactic cluster NGC 3680, error bars are indicated. Filled triangles are determinations from Keller et al. (2001): with continuous error bars amounts corresponding to the SMC cluster NGC 330, NGC 1818, NGC 2100 and with dashed error bars result for the LMC cluster NGC 2004. Open circle: determination in Hyades cluster from Lebreton et al. (2001) (upper limit for overshooting). Filled diamonds: SMC and LMC determinations performed in this work. Errors on Z have been evaluated assuming an internal error on of 0.1 dex.  
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In this paper we have estimated the overshooting distance from a turbulent core for intermediatemass main sequence stars. The result for SMC is , and for the LMC . The main contributions to errors are those brought by distance modulus and reddening uncertainties. We have shown that chemical gradients within the clouds and rotation surface effects of studied stars cannot significantly influence our results. Binary rate and IMF slope have no important effects as well. For SMC, despite different methods and data, we find a result very similar to Keller et al.'s (2001) one for cluster NGC 330. The case of LMC is more questionable because of the rather large uncertainty on reddening.
Figure 7 tends to indicate a sensitivity of overshooting to metallicity. However a mass effect cannot be excluded; we can only stress that if such a dependence exists, it should be an increase of overshooting with decreasing metallicity. However, the overshooting is expected to increase with mass, unfortunately samples studied at solar metallicity have often lower masses than those at low metallicities. Therefore further investigations are needed to disentangle these effects. In any cases, if this dependence is confirmed the next challenge will be the physical explanation of this metallicityovershooting effect.
Finally, the overshooting amounts derived in this work have a statistical meaning: they are average values over time (in real stars, "overshooting'' likely changes during the main sequence) and over mass in the considered range. Moreover these amounts represent an extramixing above the classical core generated either by inertial penetration of convective bubbles or shear phenomena related to rotation. The real extent of the core likely results from a combination of both processes; indeed, rotating models Maeder & Meynet (2001)'s rotating models still need overshooting.
Acknowledgements
We thank JeanPaul Zahn and Ian Roxburgh for helpful discussions; we are also grateful to the OGLE group for providing their data and to Pierre Morel for writing the CESAM code. We thank the referee Dr. S. C. Keller for valuable remarks and suggestions.
As we selected the data using a criterion involving the photometric standard deviation of magnitude measurements, we have to generate an artificial standard deviation for the theoretical magnitude computed from evolutionary models. Moreover the general properties of the synthetic standard deviation distribution must be similar to the OGLE 2 one.
We describe here the scheme used to generate the pseudosynthetic photometric
standard error distributions. The prefix "pseudo'' means that we have extracted information
about the standard error distribution from the OGLE 2 data themselves (see Fig. A.1a).
For that purpose, we divide the relevant range of magnitudes into bins; in each bin, we construct
the histogram of standard deviation values (Fig. A.1b). This
histogram then is fitted with a function of the form:
In our population synthesis code, for a given magnitude value m, a standard deviation value is randomly determined following the probability law derived from OGLE. After that, either the object is rejected (if the value is too large) or the magnitude m is changed into , following a gaussian distribution having a standard deviation .
Let us comment about differences between Figs. A.1a and A.1c. Figure A.1a contains the "evolutionary information''
 i.e. more objects at high magnitudes whereas Fig. A.1c does not contain this information,
objects have been uniformly distributed with respect to the magnitude. These facts explain the difference
between both figures.