A&A 488, 167-179 (2008)
DOI: 10.1051/0004-6361:20079049
G. G. Sacco1,2 - E. Franciosini1 - S. Randich4 - R. Pallavicini1,3
1 - INAF, Osservatorio Astronomico di Palermo, Piazza del Parlamento, 1, 90134 Palermo, Italy
2 - Consorzio COMETA, via S. Sofia, 64, 95123 Catania, Italy
3 - INAF, Headquarters, viale del Parco Mellini 84, 00136 Roma, Italy
4 - INAF, Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 50125 Firenze, Italy
Received 12 November 2007 / Accepted 17 May 2008
Abstract
Aims. We performed a detailed membership selection and studied the accretion properties of low-mass stars in the two apparently very similar young (1-10 Myr) clusters Ori and
Ori.
Methods. We observed 98 and 49 low-mass (0.2-1.0 )
stars in
Ori and
Ori respectively, using the multi-object optical spectrograph FLAMES at the VLT, with the high-resolution (
)
HR15N grating (6470-6790 Å). We used radial velocities, Li and H
to establish cluster membership and H
and other optical emission lines to analyze the accretion properties of members.
Results. We identified 65 and 45 members of the Ori and
Ori clusters, respectively, and discovered 16 new candidate binary systems. We also measured rotational broadening for 20 stars and estimated the mass accretion rates in 25 stars of the
Ori cluster, finding values between 10-11 and
yr-1 and in 4 stars of the
Ori cluster, finding values between 10-11 and
yr-1. Comparing our results with the infrared photometry obtained by the Spitzer satellite, we find that the fraction of stars with disks and the fraction of active disks is larger in the
Ori cluster (
% and
%) than in
Ori (
% and
%).
Conclusions. The different disk and accretion properties of the two clusters could be due either to the effect of the high-mass stars and the supernova explosion in the Ori cluster or to different ages of the cluster populations. Further observations are required to draw a definitive conclusion.
Key words: stars: formation - stars: pre-main sequence - stars: late-type - Galaxy: open clusters and associations: individual: Ori,
Ori
The time evolution of accretion properties and disk dissipation mechanisms
in young stars are two of the main open issues in star formation theory.
The analysis of the accretion properties of the stellar populations
belonging to young clusters using high-resolution (
)
spectroscopy is a powerful tool to investigate these questions.
Ori and
Ori are two of the richest young clusters
near the Sun in the age range (1-10 Myr) during which young stars
lose their circumstellar disk and stop accreting material from it
(Hartmann et al. 1998; Haisch et al. 2001; Sicilia-Aguilar et al. 2006,2005).
The Ori cluster was discovered by the ROSAT satellite
(Walter et al. 1997; Wolk 1996) around the O9.5 V binary star
Orionis AB (distance
352+166-85 pc,
Perryman et al. 1997). Because of its proximity and very low reddening
(Oliveira et al. 2004), during the last decade
Ori has become
one of the best studied young clusters. Its low-mass and substellar
population, extending down to planetary-mass objects, has been extensively
observed by photometry and low-resolution spectroscopy, both in infrared and
optical bands (Barrado y Navascués et al. 2003; Scholz & Eislöffel 2004; Sherry et al. 2004; Zapatero Osorio et al. 2002; Burningham et al. 2005; Béjar et al. 2004,2001; Zapatero Osorio et al. 2000; Kenyon et al. 2005; Caballero et al. 2007), while its high-mass
stellar content has been studied by Caballero (2007). The estimated
median age of the cluster ranges from 1 to 8 Myr, depending on the assumed
distance and measurement method adopted by different authors. Using
measurements of tangential and radial velocities, Jeffries et al. (2006)
and Caballero (2007) argued that, in the same region of the
Ori cluster, there is a sparser, kinematically separate young
stellar population belonging to the Orion OB1b association.
Zapatero Osorio et al. (2002), using low-resolution spectroscopy of a sample of 27
low-mass and substellar objects, estimated a fraction of about 30-40% of
classical T Tauri stars (CTTSs), while the fraction of circumstellar disks
(
%) was first determined by
Oliveira et al. (2006,2004), using infrared photometric data
in the K and L bands. Recently, Hernández et al. (2007) used a near-infrared
survey carried out by the Spitzer satellite to investigate the disk
properties of 336 candidate members, finding that 34% of them harbor a
circumstellar disk (27% having thick disks and 7% having evolved optically
thin disks). Moreover, they found that the total fraction of stars with
disks decreases with increasing mass from 36% (31% with thick disks) for
low-mass T Tauri stars to 15% (4% with thick disks) for Herbig Ae/Be
stars. Finally, the X-ray properties of high-mass and low-mass cluster
members have been studied by Sanz-Forcada et al. (2004) and
Franciosini et al. (2006) using the XMM-Newton satellite.
Orionis is an O8 III star (distance 400 pc,
Murdin & Penston 1977) that excites the HII region S264, delimited by a
dust ring with a diameter of 9
discovered by the IRAS
satellite (Zhang et al. 1989). The cluster, distributed over an area of 1 square degree around the O8 star, was discovered by Gomez & Lada (1998) together with two clusters located in the nearby clouds B30 and B35, by analyzing the spatial correlation among the H
sources discovered by
Duerr et al. (1982). Dolan & Mathieu (1999), using medium-resolution
spectroscopy, selected 72 members brighter than R=16 around
Orionis and found that the fraction of CTTSs (7%) belonging to
the cluster is very low if compared to other clusters and star-forming
regions in the same age range (1-10 Myr). Dolan & Mathieu (2001,2002)
extended this analysis to the whole region and suggested that the star
formation process started 8-10 Myr ago, with an accelerating star-formation
rate, and stopped 1-2 Myr ago after a supernova explosion, which shredded
the central cloud and formed the current gas ring. They also suggested that,
before the supernova explosion, the cluster was still bound to its natal
cloud, therefore the young stars closer to the central OB stars,
including the supernova progenitor, lost their circumstellar disks due
to photoevaporation by the far-ultraviolet emission from the high mass
stars. Barrado y Navascués et al. (2004) selected 170 candidate members through deep
optical photometry and performed low-resolution spectroscopy of 33 very
low-mass and substellar objects. The properties of circumstellar disks have
been studied by Barrado y Navascués et al. (2007) with the Spitzer
satellite. They found that the total fraction of members with disks (both
thick and evolved optically thin disks) is 31%, but the fraction of stars
with a thick disk is only 14%. They also found that the distribution of
Class II stars is inhomogeneous, namely, most of them are located in a
filament that goes from
Orionis to the B35 cloud. Moreover, since
several Class II stars are located near the cluster center, they argued that
high-mass stars and the supernova explosion had no effect on the
circumstellar disks.
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Figure 1:
Color-magnitude diagrams of the selected objects in the
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Although the two clusters have been extensively observed during the last
decade, only few spectroscopic data at a resolution equal to or higher than
R=10 000-15 000 are available and, especially for the Ori cluster,
most of the cluster members have been selected by means of photometry only.
Therefore, the catalogues of members of both clusters are likely contaminated by foreground field stars or young stellar objects belonging to different populations, and the accretion properties of the known members are poorly determined.
We performed high-resolution spectroscopy of two large samples of stars in
the Ori and
Ori clusters, using FLAMES at the VLT
(Pasquini et al. 2002), in order to select new high-probability
cluster members, to identify CTTSs using H
emission and other
accretion indicators and to measure mass accretion rates (MARs). Moreover,
the comparison between the ``twin'' clusters
Ori and
Ori is
a powerful tool to investigate the origin of the presumed lack of CTTSs in
the
Ori cluster. In Sacco et al. (2007) we reported a first
important result for the
Ori cluster, obtained from the analysis of
these data, namely, the discovery of three Li-depleted stars, with isochronal and nuclear ages greater than 10-15 Myr. In this paper, we focus
on the other results obtained by these data.
The paper is organized as follow: in Sect. 2 we describe target selection, observations and data analysis; results are presented in Sect. 3 and discussed in Sect. 4, where we also perform a comparison with the infrared data obtained by the Spitzer satellite. Conclusions are summarized in Sect. 5.
We selected sources in the Ori and
Ori
cluster regions in the spectral range K6-M5, which, for a 5 Myr cluster
at a distance of about 400 pc, corresponds to the magnitude ranges 13<R<18 and 11 <J< 15. All selected sources, except 4 objects without optical photometry, are
plotted in Fig. 1, where different symbols indicate the membership information available before this work.
For the Ori cluster, targets were selected using optical, infrared and X-ray data. Optical and infrared data were retrieved from the literature (Zapatero Osorio et al. 2002; Sherry et al. 2004; Wolk 1996; Kenyon et al. 2005) and the 2MASS (2 Micron All Sky Survey) catalogue
(Skrutskie et al. 2006). X-ray data were retrieved from Franciosini et al. (2006). We selected a total of 98 objects; of these, 18 are probable members on the basis of low-resolution spectroscopy, 66 are candidate members on the basis of optical and/or infrared photometry compatible with an age
10 Myr (taking also into account the error on distance) and the remaining 14 objects appear to be non-members based on photometry, but have no spectroscopic information. Of the 98 targets, 53 were also detected in X-rays.
In the Ori cluster we observed 49 sources included in the catalogues by Dolan & Mathieu (1999) (5 objects), by Barrado y Navascués et al. (2004) (34 objects) or by both (10 objects). The stars included in the Dolan & Mathieu (1999) catalogue are probable members on the basis of
medium-resolution spectroscopy, while the other stars are candidate members
based on photometric selection, with 3 objects confirmed as members by
low-resolution spectroscopy.
The sources selected in the Ori and
Ori clusters are
listed in Tables 1 and 2,
respectively. Optical and infrared magnitudes are retrieved from the
literature; spectral types for 18 stars are from previous spectroscopic
studies (Barrado y Navascués et al. 2004; Zapatero Osorio et al. 2002), while for the other stars
they have been derived from the R-I color, using the scale of
Kenyon & Hartmann (1995) interpolated for each half subtype. We did
not derive spectral types from our data, because the spectral range covered by our spectra does not include enough spectral features for the classification, and because veiling affecting
accreting objects prevents the use of spectroscopic indices based on flux
ratios. For the stars with a spectroscopic classification already performed
by other authors, we checked that the discrepancy between photometric and
spectroscopic spectral type is
1 subtype for 12 stars, between 1 and 2 subtypes for 4 stars
and of 2.5 subtypes in 2 cases (S07 and L44).
We cross-correlated our list of targets with the catalogues of stars
observed by the Spitzer satellite, published by Hernández et al. (2007) for Ori and by Barrado y Navascués et al. (2007) for
Ori. We have 83 sources in common with the Hernández et al. (2007) catalogues, 78 with the catalogue of probable members and 5 with the catalogue of uncertain members, and 44 sources in common with the Barrado y Navascués et al. (2007) catalogue. Taking into account the classification based on the slope
of the spectral energy distributions in the 3.6-8.0
m spectral range, in the
Ori cluster sample there are 34 stars harboring a circumstellar disk (
)
and 49 diskless stars, while in the
Ori sample there are 11 stars with a circumstellar disk and 33 diskless stars.
Observations were carried out using the fiber-fed multi-object spectrograph
FLAMES (Fiber Large Array Multi Element Spectrograph), mounted on the UT2
telescope at the VLT (Pasquini et al. 2002) and operated in the
MEDUSA mode (132 fibers, each with an aperture of 1.2 arcsec on the
sky). We used for both clusters the high resolution HR15N grating
(6470-6790 Å, spectral resolution R=17 000 and nominal dispersion
0.1 Å pixel-1), which includes the lithium line at 6708 Å, the
H
line (6563 Å) and other emission lines indicative of accretion
and outflow phenomena (NII at 6583 Å, HeI at 6678 Å and SII at 6716
and 6731 Å). For
Ori only, we used also the HR21 grating
(8480-9000 Å, spectral resolution R=16 200 and nominal dispersion
0.13 Å pixel-1), which includes the Ca II infrared triplet, but,
since no star in our sample shows signs of Ca II emission due to
accretion phenomena (see Sect. 3.3), these spectra have been used
only for the measurement of the radial velocity (RV). For both clusters the
FLAMES field of view (diameter 25
)
was centered around the high-mass
central star which gives the name
to the cluster (
,
and
,
,
Equinox J2000, for
Ori and
Ori, respectively).
Observations were performed in service mode and were divided into separate runs of 1 h duration each, including instrument overheads. The Ori cluster was observed in 6 runs in October and December 2004, while
Ori was observed in 8 runs in
October and November 2005. The observation log is reported in Table 3.
Table 3: Observation log.
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Figure 2:
Spectra of 4 stars of the sample. The three upper rows show spectra obtained with the HR15 grating: in the left panels we show the whole spectral range, while the middle and
right panels show the range around the H![]() |
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Data reduction was performed using the GIRAFFE girBLDRS
pipeline vers. 1.12, following the standard steps (Blecha & Simond 2004),
which include bias subtraction, division by a normalized flat-field,
correction for the differences in the fiber transmission, and wavelength calibration
using a dispersion solution from a Thorium-Argon arc lamp. The spectra
processed by the pipeline are not corrected for instrumental response and sky
background, but, except for the RVs, all measurements were performed on the sum of the spectra recorded in the different runs, after subtracting a sky background spectrum.
The sky background was calculated independently for each observing run, using a set of fibers homogeneously distributed over the field (10 for Ori and 20 for
Ori). For each run, we averaged a subsample of sky spectra without
intense spikes. To increase the signal-to-noise ratio (S/N) for line measurements, we then summed, for each source, the spectra from thedifferent runs. The S/N of the summed spectra ranges from
10 for stars with
,
to
90 for stars with
and to
170 for stars with
.
Some examples of stellar spectra are shown in Fig. 2.
We measured, independently, the RVs in each observing run, to identify cluster members and to select candidate binary systems. RVs were measured by Fourier cross-correlation (Tonry & Davis 1979), using the IRAF task FXCOR. In brief, all the spectra were cross-correlated with one template spectrum chosen among the stars belonging to the observed sample (stars S07 and S56 for
Ori and L04 and L17 for
Ori). Specifically, stars S07 and L04 have been used only for earlier type stars, and stars S56 and L17 for the others. The template spectra have been selected on the basis of their spectral type (K and M), high S/N (40-200), low rotational velocities (line full width at half maximum <0.6 Å) and because they do not show any accretion signatures. However, the H
,
other emission lines and telluric lines
have been excluded from the cross-correlated spectral range. In order to determine the RVs relative to the solar system, we measured the centroid shifts of a selected group of lines in the template spectra, using the IRAF task RVIDLINES.
For the stars with a single main peak in the cross-correlation function and with all the RVs in agreement within 2,
we computed a mean RV as the weighted average of the different measurements. The measured RVs are given in Tables 1 and 2
for the
Ori and the
Ori cluster, respectively.
The other stars, classified as candidate binaries, are discussed in Sect. 3.1.
We were able to measure the rotational velocities ()
of 20 stars and to estimate
upper limits of all the remaining members identified in the two clusters (see Sect. 3.2) and not classified as binary systems (see Sect. 3.1).
Rotational velocities were measured by cross-correlating the spectra of each star
with the spectra of the stars L04 (for stars with spectral-type earlier
than M2) and S88 (for stars with later spectral-type stars),
using the IRAF task FXCOR. Both stars used as templates show very narrow lines
(full width half maximum -
FWHM-<0.540 Å) and have S/N >20.
Specifically,
is related to the FWHM of the Gaussian that better describes the peak of the cross-correlation function by a calibration function. In order to derive the calibration function, we first cross-correlated each template spectrum with itself artificially
broadened at different velocities using the rotational profile of Gray (1992); then, we fitted the relation between the values of FWHM and the rotational velocities with a third order polynomial.
Errors on
depend on the uncertainties in the cross-correlation process between the template and the stars spectra; these have been estimated by varying
the Gaussian fit parameters within reasonable ranges. Moreover, taking into account the features of the template spectra and the uncertainties in our procedure, we conclude that the lowest measurable rotational velocity is
17 km s-1, which corresponds to a line broadening equal
to the instrumental broadening.
Inferred
and upper limits are listed in Col. 12 in Tables 1 and 2.
We measured the equivalent width of the Li line at 6708 Å and of 5 emission lines (H,
NII at 6583 Å, HeI at 6678 Å and SII at 6716 and 6731 Å), using the IRAF task SPLOT, by integrating the area under the continuum level. For the lithium line, we integrated all the spectra
over the same range (6705.8-6709.5 Å at a rest reference frame).
For the emission lines, we integrated over a spectral range different from
star to star, because accretion lines can be shifted or broadened, depending on the
velocity of the outflowing or accreting material. The lithium line measurements are reported in Tables 1 and 2, while the
emission line measurements are reported in Tables 4
and 5, for the
Ori and the
Ori clusters,
respectively. Due to the presence of molecular bands, which strongly affect the spectrum of
late-type stars, the EWs are measured with respect to a false continuum. In the rest of the paper we refer to these measurements as pseudo-equivalent widths (pEWs).
We did not measure any pEW of the double-lined binaries and of the star L48,
because in the former case we could not separate the contributions of the
different components of the binary system, while in the latter the S/N
is too low to estimate the continuum level after the sky subtraction.
For the Li and H
lines, we derived the pEWs as the
average of 3 independent measurements carried out by selecting a minimum, a
maximum and a median continuum level. Errors have been computed as the
half-difference between the maximum and minimum of the 3 measurements.
We measured the Li pEW by integrating over the 6705.8-6709.5 Å range
even in the case where the presence of the Li line was not evident. In these
cases, the line pEW is always below 150 mÅ, while in the other cases the pEW
is always higher than 250 mÅ except for the star S55, which, as discussed in Sacco et al. (2007), has started to deplete its photospheric lithium.
For the pEWs of other emission lines, we took a single measurement and, whenever the presence of the line was not evident in the spectrum, we estimated an upper limit on the basis of the faintest measurable lines observable in the same spectral range.
We measured the width of the H
line at 10% of the peak in order to discriminate between accreting and non-accreting objects and to derive the mass accretion rates (MARs), using the following relationship:
The H
widths of cluster members and MARs of the stars harboring a
disk are reported in Tables 4 and 5 for
Ori and
Ori, respectively. Errors in the H
widths
depend on the uncertainties in the continuum flux, which affect the
determination of the 10% level, while the errors in the MARs depend both on
the errors on the H
width and in the uncertainties in the
parameters included in Eq. (1). For binary systems the H
widths have
been considered undetermined in order to avoid systematic errors related to
the presence of two spectra shifted at different velocities.
We investigated the variability of the H
pEWs for all the stars, and of the H
width at 10% of the peak for the CTTSs, by comparing the measurements performed in each individual observing run. Among the WTTSs the median variability ((max(pEW)-min(pEW))/average(pEW)) ranges between 0.1 and 1.9, with a median value of 0.37, but it is larger than 1.0 only for three stars
(S28, L06, L36). For the latter stars, the variation is observed only in a single run, suggesting that it was probably recorded during a flaring event.
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Figure 3:
H![]() ![]() ![]() |
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The variability of the pEWs and of the 10% widths for the CTTSs
is plotted in Fig. 3 as a function of the 10% widths measured
on the summed spectra. As shown in the figure, we do not find evidence for
any correlation between the two measurements of variability and the 10% width,
which, as explained in the previous section, is correlated with the mass
accretion rate. The median value, the mean and the standard deviation of
the variability of the pEWs are 0.50, 0.53 and 0.31, respectively, while the
median, mean and standard deviation of the 10% widths are 0.22, 0.26 and
0.15, respectively. In 5 cases (S05, S33, S65, S98
and L34) for the pEWs and in 4 cases (S12, S65, S69 and L34) for the 10% widths, the variations differ by more than 3
from the mean of the sample. In the case of the star S23, which is not included in Fig. 3,
we find a variability of the H
pEW
4, because the line is composed of both absorption and emission features, and the pEW is positive in some runs
and negative in others.
Moreover, the intrinsic variability is in most cases higher than
the errors in the measurements of the H
pEWs and 10% widths, which
are of the order of 10%, as can be seen from Tables 4 and 5.
We found 11 stars in Ori and 5 in
Ori for which
at least 2 of the measured RVs differ by more than 2
,
and therefore we classified them as candidate binaries. The measured RVs in each run are given in Tables 6 and 7. Specifically, in the
Ori sample, for 6 binaries the RVs are not in agreement between each other. Of these, 5 stars show a cross-correlation function with two distinct peaks, corresponding to the RVs of both components of the binary system.
The RVs as a function of time for these 6 binaries are plotted in Fig. 4, and show a
well-defined sinusoidal trend. We have performed a least-squares fit to the RV curves using a sine
function; in the case of double-lined binaries, both curves were fitted
simultaneously.
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Figure 4: Velocity vs. time for the 6 binary systems reported in Table 8. In the left and right panels we show the values obtained from the runs performed in October and December, respectively. The continuous and dotted curves show the sine functions that fit the data. |
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The binary system parameters derived from the fits with their 1
uncertainties are given in Table 8. For the other 5 candidate binaries in
Ori, only the RV measured in October 2004 does not agree with the other 5 RVs measured in December 2004, that are compatible within the errors. A careful check of the October spectra did not reveal any reduction or analysis problem, therefore we believe that the observed discrepancy is real. These stars might therefore be binaries with a period of some months, for which RV variations cannot be detected in the observations taken in December, which are distributed over a period shorter than 10 days. Further RV measurements are required to confirm their nature.
In the
Ori cluster, we found one star with two distinct peaks in the cross-correlation function, for which however we cannot obtain a fit to the RV data, and 4 stars with at least two of the measured RVs differing by more than twice the error bar.
For the selection of members from RVs, we used the derived center of mass velocities for the binary systems listed in Table 8. For the other binaries, the system velocities cannot be determined, and will therefore not be used for their membership determination.
To determine cluster membership we used three independent criteria, based on the RV distribution, the pEW of the Li absorption line, and the presence of the H
line in emission.
In Fig. 5 we show the RV distributions for the observed
stars in Ori and
Ori, excluding the 10 probable binaries
(5 for each cluster) for which the velocity of the center of mass is
undetermined. The two distributions have been fitted with the weighted sum
of two Gaussian functions (shown in Fig. 5 with dashed
curves), one describing the velocity distribution of the cluster members and
the other describing the velocity distribution of the field stars. For
Ori the cluster distribution is centered at
km s-1, with a standard deviation
km s-1, while for the field star distribution we find
km s-1 and
km s-1. In
the case of
Ori we obtain
km s-1 with
km s-1 for the cluster and
km s-1 with
km s-1 for the field,
although, considering the low number of field stars in the
Ori
sample, the parameters of the latter distribution are poorly determined. For
each cluster, we classify a star as a member if its RV, taking the error bar
into account, differs by less than 3
from the cluster
average velocity
.
Using this criterion, among the 93 stars in
the
Ori sample we find 64 stars with RV consistent with membership,
while among the 44 stars in the
Ori sample, 37 have RV consistent
with membership. The membership of the 10 binaries excluded from the RV distributions is undetermined. The expected contamination of the
Ori member sample by non-members, estimated by integrating the field star distribution between
and
,
is
2 stars.
The velocity of the Ori cluster is in agreement with that
of its central high-mass star
Orionis (
km s-1,
Morrell & Levato 1991) and with the cluster mean velocity found by
Zapatero Osorio et al. (2002), Kenyon et al. (2005) and Jeffries et al. (2006).
Our measurements do not suggest the presence of the kinematically-separated
(
km s-1) young stellar population discovered by
Jeffries et al. (2006). Since this second population is concentrated to the
north-west of the hot star, this apparent disagreement can be attributed to
the different positions of the observed fields. In fact, if we limit the
analysis to the Jeffries et al. (2006) field with the largest area in
common with our field, and compute the number of stars in the two RV ranges
(20-27 and 27-35 km s-1) defined by Jeffries et al. (2006) for the two
populations, we find in our field a ratio (3/65 or
%) consistent
with that found by Jeffries et al. (2006) (4/46 or
%). The
velocity of the
Ori cluster is in agreement with that found by
Dolan & Mathieu (1999) (
km s-1), but slightly lower than
the velocity of its high-mass central star
Orionis (
km s-1, Kharchenko et al. 2007).
Late-type stars (0.08-0.5 )
deplete their photospheric lithium
during the pre-main sequence (PMS) phase (Siess et al. 2000; Bodenheimer 1965), therefore the
presence of strong Li absorption at 6708 Å in young clusters can be
used as an independent criterion to identify cluster members. Figure 6 shows the Li
pEWs measured in the two clusters as a function of the R-I color.
Since in all cases where the Li line was not clearly identified the
Li pEW turned out to be less than 250 mÅ, we fixed the threshold for
selecting cluster members at this value. For the double-lined binaries and
for the star L48, for which we cannot measure the Li pEW, we
established the membership only by considering whether the Li absorption
feature can be identified or not.
As shown in Fig. 6, the bulk of the measured pEWs above the
chosen threshold is located around a median value of 560 mÅ with a
dispersion of
100 mÅ. According to the curves of growth
derived by Palla et al. (2007) for a sample of stars observed by GIRAFFE and very
similar to that described in this paper, this median pEW is compatible with
the full preservation of lithium. In the
Ori cluster, we found a
few sources with pEWs between 250 and 400 mÅ. As shown by Sacco et al. (2007), these low pEWs are mainly due to spectral veiling, but we cannot exclude the presence of some partially-depleted stars. This hypothesis is supported by the discovery of three highly-depleted stars among the
Ori cluster members, reported by Sacco et al. (2007).
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Figure 5: Radial velocity distribution of the observed sources in the two clusters. Only single stars and the binaries of Table 8 are included. The dashed curves are the functions resulting from the fit of the velocity distribution with a weighted sum of two Gaussians. |
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Table 8:
Binary systems parameters for the 6 binaries in the Ori cluster for which a sinusoidal fit of the RV curves was performed. Errors are
.
Young K and M stars show emission in H
due to accretion and chromospheric activity (Barrado y Navascués & Martín 2003). Therefore we used the presence of H
in emission as the third criterion to identify cluster members. Although, due to the very low S/N, we have not
been able to subtract the correct emission of the sky from the spectrum of the star
L48 (see Sect. 2.5), this star has been classified as a member, considering
that the H
emission line observed in its spectrum is nearly two times larger than the H
line observed in the sky spectrum.
In Tables 1 and 2,
for the Ori and
Ori clusters respectively, we give
the results of our membership selection from the individual criteria
mentioned above, together with a final membership assessment derived
from the comparison of the three criteria, which in most cases agree.
In Table 1 we also indicate if the X-ray counterparts of
the stars have been detected by XMM-Newton.
For the probable binaries with undetermined RV, we relied only on the Li and
H
criteria, considering as possible members those stars for which
both Li and H
indicate membership, although we cannot exclude their
belonging to a separate young population present in the same area.
In the
Ori sample, we classify 62 stars as members or possible
members and 29 stars as non-members. In
Ori, we find 42 members or
possible members and 2 non-members. For the remaining 7 stars in
Ori and 5 stars in
Ori, the three criteria do not agree,
and we assigned a final membership based on the following considerations:
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Figure 6:
Li pEWs vs. R-I for the stars of the ![]() ![]() |
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Figure 7:
CMD of the observed targets in the ![]() ![]() ![]() |
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Figure 8:
H![]() ![]() ![]() ![]() |
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We have compared our results with the previous membership information
available in the literature (see Sect. 2.1). In the case of
Ori, among the 18 members selected on the basis of previous
spectroscopic information, we confirm membership for all of them, except for
the star S37, which has been classified as a non-member because of its RV (see
above). Of the 66 candidates selected only by means of the photometric
data, only 47 are confirmed spectroscopically as members, while the
remaining 19 are non-members. For the remaining 14 targets that
appear to be non-members on the basis of photometry, we confirm that
13 of them are indeed not members of the cluster. The remaining star (S47)
has been classified by us as a member, because it has all the spectroscopic
indicators consistent with membership, and is also an X-ray source.
However its position in the CMD is anomalous, falling below the ZAMS
(see Fig. 7). Given the spectroscopic and X-ray evidence,
we believe that its photometric data, retrieved from Wolk (1996),
might be wrong, therefore its photometry will not be taken into account in the following sections.
For Ori, we find that all the targets selected by means of
previous spectroscopy from Dolan & Mathieu (1999) (15 stars) and
Barrado y Navascués et al. (2004) (3 stars) are confirmed as members, while among the
candidates selected only by means of photometry we find 26 members or
probable members and 4 non-members.
In Tables 4 and 5 we report, for
Ori and
Ori respectively, all data concerning accretion,
outflow and disk signatures, namely, the H
width at 10% of the
peak, the pEWs of H
and the pEWs of the HeI emission lines,
generally interpreted as due to the accretion flow, the pEWs of the forbidden NII
and SII lines at 6583, 6716 and 6731 Å, which are
signatures of outflow material and, in the last column, the
Spitzer classification performed by Hernández et al. (2007)
and Barrado y Navascués et al. (2007), using the criterion defined by Lada et al. (2006)
based on the slope
of the spectral energy distribution
between 3.6 and 8.0
m, where
.
In Tables 4 and 5 we also give the accretion
rates derived from Eq. (1) for each star harboring a circumstellar disk
with an H
width at 10% of the peak larger than 200 km s-1.
In the
Ori cluster, MARs range between
yr-1 and
yr-1, while in the
Ori cluster they range between
yr-1 and
yr-1.
The H
pEWs as a function of the H
widths at 10% of
the peak are shown in Fig. 8, where stars harboring
a circumstellar disk are indicated by filled symbols. As can be seen in this figure and in Tables 4 and 5, we find some stars with an H
larger than
200 km s-1 without any other accretion signature and classified as class III objects by
the Spitzer data. This result is likely due to the
presence of absorption bands in the H
spectral region, which causes
an overestimate of the H
width, and to the rotational broadening (the stars L16 and L33, which are the most relevant cases have
km s-1).
We also find some stars with H
higher than 200 km s-1 but
with a low pEW (S23 and S47 is the most relevant case), and stars in the
opposite situation (S93 and S94 are the most relevant cases). The former result is mainly due to the presence of intense absorption features overlapping
the emission ones as already pointed out by Sicilia-Aguilar et al. (2006)
for some CTTSs belonging to the Tr37 cluster, while the latter result is due
to the presence of accretion flows only in the direction perpendicular to the
line of sight.
We also observed the Ori sources in the spectral range including
the CaII infrared triplet, which can also be used to measure MARs (Mohanty et al. 2005), but in the whole sample we do not find any star clearly showing CaII emission due to accretion.
As mentioned in Sect. 2.1 in the Ori sample there
are 78 sources in common with the Hernández et al. (2007) catalogue of cluster
members and 5 sources in common with the catalogue of uncertain members. According to our membership
analysis, in the former catalogue there are 62 members and 16 non-members, while in the latter one there are 3 members and 2 non-members. Given that 20% of the stars in common with our catalogue
are non-members, the number of
Ori members included in the
Hernández et al. (2007) catalogue, selected mainly using photometric data,
is likely overestimated, and, therefore,
the fraction of T Tauri stars (0.1-1.0
)
with disks
(
)
derived by them is likely underestimated. In the
Ori cluster, we have 44 sources in common with the Barrado y Navascués et al. (2007) catalogue, 40 classified as members
and 4 as non-members according to our analysis.
According to the disk classification by Hernández et al. (2007) and Barrado y Navascués et al. (2007), among the Ori cluster members we have 1 protostar, 28 stars with thick disks, 5 stars with evolved disks and 31 diskless stars, while in
Ori we have 7 stars with thick disks, 4 with evolved disks and 29 diskless stars. Therefore, in the mass range 0.2-1.0
,
the fraction of stars with disks
in the
Ori cluster (
%) is larger than in the
Ori
one (
).
The discrepancy between the two clusters does not depend on the analysis method, because we compare the two clusters using homogeneous data and the same selection criteria both for membership and disk classification. We can also exclude that this result might be due to a bias in the target selection against the diskless stars. In fact, for both clusters the observed stars have been retrieved from catalogues of candidate members selected through CMDs and spectroscopic membership indicators (lithium and RVs), which do not depend on disk and accretion properties.
Another possible bias in the target selection can be present against the
older stars, which are more likely class III objects. However, in contrast
with our results, this kind of bias should affect more strongly the
Ori sample than the
Ori one. In fact, as pointed out in
Sects. 2.1 and 3.2, the
Ori targets are selected mainly from the Barrado y Navascués et al. (2004) catalogue which includes only stars around the 5 Myr isochrone (Baraffe et al. 1998, models at 400 pc),
while the
Ori target sample includes stars in a larger age range
(0-20 Myr). Moreover, the agreement between the fraction of stars with disks
in our sample (
)
and in the sample of all stars in the magnitude range
11.3<J<14.8 (105 stars) included in the Barrado y Navascués et al. (2007) catalogue (
%)
confirms that our sample is not affected by target selection biases.
On the other hand, the disagreement between the fraction of stars with disks
found in the
Ori sample (
%) and that obtained, in a
very similar magnitude range (
11.5<J<14.6) by Hernández et al. (2007)
confirms that, as we argued above, their sample of cluster members
is contaminated by a large number of field stars.
The fraction of stars with disks found by us in Ori is in
agreement with that found by Caballero et al. (2007) for the brown dwarfs
(
%). This suggests that the disk frequency does not depend on mass
over the interval 0.02-1.0
.
Table 9:
Disk and accretion properties a in
Ori and
Ori
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Figure 9:
Age distribution of the members selected in ![]() ![]() |
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In Table 9 we compare the disk and accretion properties
of Ori and
Ori on the basis of all the different disk and
accretion classifications. Furthermore, in the last two columns, we
compare the fraction of active disks, namely, the fraction of accreting
objects among the stars harboring a circumstellar disk (class I, class II
and EV).
The data reported in Table 9, together with the
comparison between the accretion rates measured by means of the Hwidths at 10% of the peak, prove that, in spite of the similarities between
the two clusters,
the disk and accretion properties of their low-mass stellar populations are
very different. Specifically, in the
Ori cluster we find a larger
fraction of stars with disks, a larger fraction of accreting objects, and
higher accretion rates than in
Ori. Furthermore, the last two
columns of Table 9 show that the two clusters differ also if we consider only the subsample of stars with a circumstellar disk.
Dolan & Mathieu (2001,1999) already described the lack of
strong H
emitters in the
Ori cluster and the existence
of a discrepancy with the fraction of CTTSs observed in the
B30 and B35 clouds, located 2.2
and 2.7
from the central
star
Orionis. They supposed that, as observed in the Trapezium around
Ori (Johnstone et al. 1998), circumstellar disks could have been photo-evaporated by the far-UV radiation of the high-mass stars in the period before the supernova explosion, when high and low-mass stars were confined by the parent cloud in a smaller region.
The same effect might have not affected the
Ori stars because its
central high-mass star is less bright than
Orionis, and likely
also than the supernova progenitor and because the
Ori cluster might never
have been confined in a smaller region as Dolan & Mathieu (2001,1999)
supposed for the
Ori cluster.
However, recent N-body simulations showed that the effect of photoevaporation of the circumstellar disks due to emission from the high-mass stars is negligible for clusters composed of fewer than 1000 members (Adams et al. 2006), while Barrado y Navascués et al. (2004) found only
170 candidate members in their survey over an area of 0.3 deg2 (with a completeness limit of
0.025 ),
including most of the
Ori cluster.
Moreover, Barrado y Navascués et al. (2007) found several stars with disks near the
central O stars and no correlation between the fraction of stars with disks
and the distance from
Orionis.
Furthermore the difference in the accretion properties of stars with disks in the two clusters, discussed in Sect. 4.2, does not support the Dolan & Mathieu (2001) hypothesis, because the far-UV radiation of high-mass stars could trigger the disk photoevaporation but should not directly influence the accretion on the stars. On the contrary, Fatuzzo et al. (2006) suggested that a supernova explosion could increase the MARs of the nearby stars. Specifically, they hypothesized that the enhancement of the cosmic ray flux following a supernova explosion causes an increase of the ionization levels in the circumstellar disks of the nearby stars and, therefore, of the magneto-rotational instabilities, which, according to the Gammie (1996) model, trigger accretion on the CTTSs
Another possibility is that the discrepancy between the two clusters could
be due to a different age. As shown by Hernández et al. (2007) using Spitzer photometry and by Haisch et al. (2001) using J, H, K and L band photometry, the
fraction of stars with disks decreases from 100% in the youngest regions to a few
percent in 6-7 Myr old clusters. Moreover, Hernández et al. (2007) and Barrado y Navascués et al. (2007) found a larger fraction of evolved thin disks in the older young clusters. Since a
significant fraction of these evolved disks might have stopped the accretion
onto the stars (Sicilia-Aguilar et al. 2006, and references therein), it is
likely that in young clusters the fraction of non-accreting disks grows with
age. Following these considerations, the suggestion that the Ori
population is more evolved than the
Ori one could well explain the differences
between the two clusters.
To investigate this hypothesis, we plot in Fig. 9 the age
distributions of the members selected in both clusters. Ages have been calculated from the Siess et al. (2000) evolutionary models using the I and R magnitudes, while for the distance modulus and the reddening we used the
same values as in Fig. 1. The two histograms clearly show that the Ori population is older than the
Ori one. Moreover, the CMD for the two clusters, shown in Fig. 7,
shows that the age difference especially concerns the lowest mass stars.
Although Fig. 9 seems to confirm the age hypothesis,
we note that the age distributions shown in Fig. 9 could
be affected by
large sources of errors: the reddening due to the presence of the
circumstellar disks could cause an underestimate of the age, while veiling
due to accretion could cause an overestimate of age; errors on distances,
which are poorly determined
in both clusters, could strongly affect the age distributions and, as shown by
Hillenbrand et al. (2008), different evolutionary models give different
results. The presence of some CTTSs among the oldest and lowest stars in Ori
confirms that ages derived from R and I magnitudes are affected by these
sources of errors, in fact errors due to accretion veiling strongly affect
the low-mass stars (Herczeg & Hillenbrand 2008).
Moreover, taking into account that, according to the Siess et al. (2000)
models, stars in the mass range 0.3-0.2
very rapidly deplete
their photospheric lithium after 15-30 Myr, the oldest stars of the
Ori cluster in this mass range should have depleted part of their lithium, but, as shown
in Fig. 6, the lithium pEWs of
Ori members are
never below 450 mÅ.
More detailed HR diagrams based on spectral types derived from low-resolution spectroscopy are required to better investigate the age difference between the two clusters.
In this paper we have reported the results of FLAMES/VLT optical spectroscopic observations of 147 low-mass stars, in the 0.2-1.0
mass range, belonging to the two very similar clusters
Ori and
Ori.
Using RVs, H
and Li line pEWs, we identified 65 bona-fide members
of the
Ori cluster and 45 members of
Ori. Furthermore
we discovered 16 new binary systems and binary candidates, 10 of which are probable members of the clusters and measured rotational velocities of
20 stars. To study the accretion properties, we estimated the stellar
MARs from the width at 10% of the peak of the H
line
and measured the pEWs of the H
and other emission lines, which
are signatures of accretion/outflow phenomena.
We compared our results with the Spitzer observations of the two
clusters by Hernández et al. (2007) and Barrado y Navascués et al. (2007), finding that:
a) the fraction of stars with disks obtained by Hernández et al. (2007)
is likely underestimated due to the presence of a large number of field stars
in their catalogue of members; b) the fraction of stars with a circumstellar
disk in the
Ori cluster (
%) is larger than in
Ori (
%); c)
the fraction of active disks in
Ori (
%) is
larger than in
Ori (
%).
We have discussed two hypotheses that could explain the discrepancy between
the two clusters: either the circumstellar disks in the
Ori
cluster dissipate more rapidly due to the effect of the massive stars
emission or the
Ori cluster could be older and more evolved than
Ori. The former hypothesis is in contradiction to some theoretical and
observational studies, while the latter one cannot be confirmed due to the
uncertainties in stellar ages. Further low-resolution spectroscopic
observations are required to reach a firm conclusion.
Acknowledgements
We thank the anonymous referee and Antonella Natta for useful suggestions and comments. We acknowledge the staff of the ESO Data Management and Operations department, who performed our observations in service mode.
This work makes use of results produced by the PI2S2 Project managed by the Consorzio COMETA, a project co-funded by the Italian Ministry of University and Research (MIUR) within the Piano Operativo Nazionale ``Ricerca Scientifica, Sviluppo Tecnologico, Alta Formazione'' (PON 2000-2006). More information is available at http://www.pi2s2.it and http://www.consorzio-cometa.it.
This research has been supported by an INAF grant on Stellar clusters as probes of star formation and early stellar evolution (PI: F. Palla). This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.
Table 1:
Photometry and membership for the observed sample stars in the Ori cluster.
Table 2:
Photometry and membership for the observed sample stars in the Ori cluster.
Table 4:
Accretion properties of the
Ori cluster members
Table 5:
Accretion properties of the
Ori cluster members.
Table 6:
Radial velocities of the candidate binary systems identified in the
Ori cluster sample.
Table 7:
Radial velocities of the candidate binary systems identified in the
Ori cluster sample.