Issue |
A&A
Volume 509, January 2010
|
|
---|---|---|
Article Number | A102 | |
Number of page(s) | 10 | |
Section | Galactic structure, stellar clusters, and populations | |
DOI | https://doi.org/10.1051/0004-6361/200913258 | |
Published online | 26 January 2010 |
The metallicity of the open cluster
Tombaugh 2![[*]](/icons/foot_motif.png)
S. Villanova1 - S. Randich2 - D. Geisler1 - G. Carraro3,4 - E. Costa5
1 - Universidad de Concepción, Departamento de Astronomia, Casilla
160-C, Concepción, Chile
2 - INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125
Firenze, Italy
3 - ESO, Alonso de Cordova 3107, Vitacura, Santiago, Chile
4 - Dipartimento di Astronomia, Universitá di Padova, Vicolo
Osservatorio 3, 35122, Padova, Italy
5 - Departamento de Astronomía, Universidad de Chile, Casilla 36-D,
Santiago, Chile
Received 7 September 2009 / Accepted 5 November 2009
Abstract
Context. Open clusters are excellent tracers of the
structure, kinematics, and chemical evolution of the disk and a wealth
of information can be derived from the spectra of their constituent
stars.
Aims. We investigate the nature of the chemical
composition of the outer disk open cluster Tombaugh 2. This has been
suggested to be a member of the GASS/Mon substructure, and a recent
study by Frinchaboy et al. (2008) suggested that this was a
unique open cluster in possessing an intrinsic metal abundance
dispersion. We aim to investigate such claims.
Methods. High resolution VLT+GIRAFFE spectra in the
optical are obtained and analyzed for a number of stars in the
Tombaugh 2 field, together with independent UBVI photometry.
Radial velocities and position in the color-magnitude diagram are used
to assess cluster membership. The spectra, together with input
atmospheric parameters and model atmospheres, are used to determine
detailed chemical abundances for a variety of elements in
13 members with good spectra.
Results. We find the mean metallicity to be
[Fe/H] = -0.31
0.02 with no evidence for an intrinsic abundance dispersion,
in contrary to the recent results of Frinchaboy
et al. (2008, MNRAS, 391, 39). We find Ca and Ba to be
slightly enhanced, while Ni and Sc are solar. The r-process
element Eu was found to be enhanced, giving an average
[Eu/Ba] = +0.17. The Li abundance decreases with
on the upper giant branch and maintains a low level for red clump
stars. The mean metallicity we derive agrees well with that expected
from the radial abundance gradient in the disk for a cluster at its
Galactocentric distance.
Conclusions. Tombaugh 2 is found to have abundances
as expected from its Galactocentric distance and no evidence for any
intrinsic metallicity dispersion. The surprising result found by
Frinchaboy et al. (2008), which is the presence of two
distinct abundance groups within the cluster, implying either a
completely unique open cluster with an intrinsic metallicity spread or
a very unlikely superposition of a cold stellar stream and a very
distant open cluster, is not supported by our new result.
Key words: Galaxy: disk - open clusters and associations: general - open clusters and associations: individual: Tombaugh 2 - Galaxy: structure
1 Introduction
One of the many regions in our Galaxy for which we lack detailed knowledge is the outer disk. How and when was it formed? By an outside-in or inside-out process (Chiappini et al. 2001)? What is the nature of the metallicity gradient? Does it show a constant slope with Galactocentric distance, or is there a leveling out in the outer disk (Carraro et al. 2007; Magrini et al. 2009; Twarog et al. 1997; Carraro et al. 2004)? What if any is the role of mergers?
A particularly intriguing development in recent years has been the suggestion that there is a merger remnant lying near the Galactic plane. The structure was first identified as the Monoceros stream (Mon; Newberg et al. 2002; Yanny et al. 2003; Ibata et al. 2003), also known as the Galactic anticenter stellar structure (GASS; Crane et al. 2003; Rocha-Pinto et al. 2003). GASS/Mon was discovered as an overdensity of stars near the Galactic plane that seems to wrap around the outer parts of the Galactic disk and is frequently explained as tidal debris from the disruption of a dwarf galaxy on a low inclination orbit (e.g., Crane et al. 2003; Peñarrubia et al. 2005; Yanny et al. 2003).
Further arguments have been made that the reputed dwarf galaxy
in Canis Major (CMa; Bellazzini
et al. 2004)
is the progenitor of the Mon/GASS structure (Martin et al. 2004).
However, the nature and/or reality of the proposed Mon ``tidal debris
stream'' and the CMa overdensity have been called into question and are
currently a matter of great debate. Momany (2006,2004) have argued that much of the
observed stellar overdensity associated with Mon - and
particularly all of that associated with
CMa (at
) - is due to the
warping and flaring of the Galactic disk, and that no
``extra-Galactic'' component is needed to account for the apparent
overdensities in the third quadrant. The presence of ``blue plume''
stars in this part of the sky has been used to argue further for the
presence of a dwarf galaxy nucleus in CMa (Bellazzini et al. 2004; Dinescu
et al. 2005; Butler et al. 2007; Martínez-Delgado
et al. 2005).
However, these young stars have also been more prosaically attributed
to the presence of previously unknown features of spiral arm structure (Moitinho
et al. 2006; Carraro et al. 2005).
Open clusters have traditionally played a critical role in the
study of the structure, kinematics, formation and chemical evolution of
the disk. For this reason alone, the old outer Galactic disk open
cluster Tombaugh 2 (To2; at Galactic coordinates [l,b] =
[232.8, -6.9]),
located at a Galactocentric distance of
15 kpc and 700 pc below the
Galactic plane, is worthy of study and was indeed the subject of
several previous metallicity investigations. Brown
et al. (1996) obtained high resolution spectra of
three stars, finding [Fe/H] = -0.4
0.25, while Friel et al. (2002)
derived [Fe/H]-0.44
0.09 from much lower resolution spectra of 12 members. Both of
these studies were conducted with the CTIO 4 m telescope.
Subsequently, Frinchaboy
et al. (2004) suggested that To2 may
possibly be associated with Mon/GASS as well as with CMa (Bellazzini et al. 2004).
Note that its Galactic longitude is very close to that of the
CMa overdensity. This provided a strong additional incentive
for a followup study of its chemistry and kinematics in greater detail
with a larger telescope, as even the giants are faint due to
its distance of over 13 kpc. Frinchaboy et al. (2008
- hereafter F08) used this motivation to obtain VLT/FLAMES spectra
(both with UVES at high resolution and with GIRAFFE at somewhat lower
resolution). They were able to derive velocities and detailed
abundances of a number of elements for 18 To2 cluster
members. Their most surprising result was an apparent large spread in
metallicity:
[Fe/H] >
0.2. They were unable to account for this spread given their
observational errors and presented a number of possible scenarios,
including the likelihood that To2 possessed an intrinsic metallicity
spread. They argued for the possible presence of two populations in the
cluster, distinguished by their different mean chemical
characteristics - with a metal-rich, Ti-normal group and a
metal-poor, Ti-enhanced group, namely (
[Fe/H]
,
[Ti/Fe]
)
(-0.06, +0.02) and (-0.28, +0.36). The more
metal-poor group appeared more centrally concentrated, and they
suggested that this group represented the true To2 clusters
stars, and the metal-rich population was an overlapping, and
kinematically associated, but ``cold'' stellar stream.
If true, this would be the first such metallicity spread uncovered in an open cluster, and/or the first evidence of multiple populations in a cluster of this age (about 2 Gyr) in the Galaxy, making To2 an extremely intriguing object. Although multiple populations within globular clusters are now known to exist, these clusters are all much more massive than an open cluster such as To2, as expected if this phenomenon is due to a cluster being massive enough to retain ejecta from a first generation of stars in order to make a second, chemically enriched generation. To date, the F08 study is the only one which suggests multiple populations in an open cluster. However, as F08 pointed out, further observations are desperately required to corroborate their surprising results and make a more definitive investigation.
Given its above history, a new metallicity study of To2 is imperative. These points motivated the present study, wherein we wished to obtain independent high-resolution spectra with a large telescope of a number of To2 stars in order to investigate the reality of the putative metallicity dispersion and the nature of its chemical composition. In Sect. 2 we describe our observations and reduction procedures. In Sect. 3 we discuss the determination of radial velocities and membership for our observed sample. In Sect. 4 we present the details of our abundance analysis, including the derivation of the input atmospheric parameters and of the internal errors. In Sect. 5 our basic abundance results are presented, and in Sect. 6 we compare these to previous investigations, especially that of F08 in view of the above. Finally, we summarize and emphasize the importance of our main results in Sect. 7.
![]() |
Figure 1: Left panel: the CMD of Tombaugh 2 in the optical and infrared. Members stars studied in this paper are indicated as filled circles. Right panel: distribution on the sky of our targets (filled points) and F08 objects (empty squares). |
Open with DEXTER |
Table 1: Target stars and parameters.
2 Observations and data reduction
Our data-set consists of high resolution spectra collected with
FLAMES-GIRAFFE/VLT@UT2 (Pasquini
et al. 2002) in Service mode from March 6
to March 25 2006, within a project devoted to measure
radial velocities, membership, and chemistry in a large sample of open
clusters (Randich et al. 2005).
The GIRAFFE spectrograph was used in the HR15N setting, providing a
resolution of
and covering a spectral range of
320 Å with the central wavelength at
6650 Å. Typical seeing during the observations was in the
range of 0.8-1.2 arcsec.
The cluster was observed with two different configurations (A
and B), centered at the same position
(RA(2000) = 07h 03m 01.95s, Dec =
-20d 49m 50.2s). We obtained four and three
45 min long exposures for the configurations A
and B, respectively. Medusa fibers were allocated to 93 and
120 stars in the two configurations, with 78 stars in
common. Hence, we obtained in total
spectra of 135 cluster candidates. These cover the magnitude
range of and
are located in different regions of the cluster color-magnitude diagram
(CMD): namely, the
turn-off and subgiant branch, the red giant branch (RGB), the red clump
(RC), and the blue plume.
In this paper we focus on RGB and RC stars for a total of
37 objects. 15 of them turned out to be members
(see below for the membership determination), but only
13 had spectra with sufficient quality in order to measure
chemical abundances and are listed in Table 1. Their position in the
cluster CMDs is shown in Fig. 1.
Targets were originally selected from the Phelps
et al. (1994) photometry. But in this paper
we use our new optical photometry (see following sub-section)
and 2MASS (JHK - Skrutskie
et al. 2006).
Data were reduced using GIRAFFE pipelines (Ballester
et al. 2000), including bias subtraction, flat-field
correction, and wavelength calibration. Sky subtraction was performed
plate by plate using the median sky as obtained from the fibers pointed
on empty regions of the field. Radial velocities were derived from each
single spectrum (see below). Spectra of stars that were
not found to be RV variables were then
co-added. The typical signal to noise ratio per pixel is .
2.1 Photometric material
We complemented the spectroscopic observations with new UBVI
observations secured with the Cerro Tololo Inter-American Observatory
1.0 m telescope, operated by the SMARTS consortium as a part
of a large photometric program to study stellar fields in the third
Galactic Quadrant.
Tombaugh 2 was observed in December 2008 under
photometric conditions. Full details of the observations, data
reduction and photometric calibration can be found in (Carraro & Costa 2009) and
references therein. Here we would just like to point out the large
field (20
on a side)
and good scale (0.297
/pixel)
provided by our set-up, which allows for an optimum study of both the
central cluster and the surrounding field. For the purpose of this
study, we selected stars within 3 arcmin from the nominal
cluster center and cross-correlated our data with the
2MASS catalog. Typical deep exposures times in U,
B, V and
were 2000, 1500, 1200, and 1200 s, respectively.
3 Radial velocities and membership
In the present work, radial velocities, coupled with the position of
the star along the principal giant sequences in the cluster
color-magnitude diagram, were used as the membership criterion
since cluster stars all have similar motion with respect to the
observer. The radial velocities of the stars were measured using the
IRAF FXCOR task, which cross-correlates the object spectrum with a
template. As a template, we used a synthetic spectrum obtained
through the spectral synthesis code
SPECTRUM, using a Kurucz model
atmosphere (Kurucz 1992) with
roughly the mean atmospheric parameters of our stars
K,
,
km s-1,
.
At the end, each radial velocity was corrected to the
heliocentric system. We calculated a first approximation mean velocity
and the rms (
)
of the velocity distribution. Stars showing
more than
from the mean value were considered probable field objects and
rejected, leaving us with 15 objects as probable members.
Only 13 of them had spectra with sufficient quality in order
to measure chemical abundances, and their positions in the CMD and on
the sky are shown in Fig. 1.
Radial velocities of these 13 targets are reported in Table 1, while coordinates, V
magnitudes and velocities for non-members and the two member
with bad spectra are reported in Table 5. We found for the
cluster a mean heliocentric radial velocity and an observed
dispersion of:

Both of these values agree excellently with F08.
Our analysis shows that only 15 out of the 37 sample
stars are cluster members, thus yielding a contamination of %.
![]() |
Figure 2:
Example of the spectral synthesis method applied to stars #236
and #591. Observed spectra are plotted with continuous lines,
while the synthetic ones are plotted with dashed lines. Abundances (
|
Open with DEXTER |
4 Abundance analysis
4.1 The linelist
The linelist for the elements we measured (with the exception of Li,
Ba, and Eu) were obtained from the VALD database (Kupka et al. 1999) and
calibrated using the solar-inverse technique. For this purpose we used
the high resolution, high S/N
Solar spectrum obtained at NOAO (National Optical Astronomy
Observatory, Kurucz et al. 1984). The EWs for the
reference solar spectrum were obtained in the same way as the observed
spectra (see next subsection) with the exception of the strongest
lines, where a Voigt profile integration was used. Lines
affected by blends were rejected
from the final line-list. Abundances were determined using MOOG, coupled with a solar
model atmosphere interpolated from the Kurucz
(1992) grid using the canonical atmospheric parameters for
the Sun:
K,
,
km s-1
and
.
In the calibration procedure, we adjusted the value of the line
strength
of each spectral line so that the abundances obtained from all the
lines of the same element yield the same value.
Note that this procedure is identical to that used in F08.
For Ba, known to be affected by hyperfine structure and
isotopic effects, we applied the spectrum synthesis method instead. For
the Sun we considered Ba lines at 4554, 5853, 6141,
and 6496 Å. Hyperfine components of those lines and isotopic
composition were taken from McWilliam (1998). For each line we
determined a Ba abundance, and then we adjusted the value
of each one in order to yield the same abundance obtained from all the
lines. For the present data we used only the Ba feature
at 6496 Å.
The Eu line we used is located at 6645 Å and it is
blended with a Si feature. We verified that at the
temperature of our stars the contribution of the Si line is
negligible. Anyway we obtain Eu abundance by spectral
synthesis method. For Si abundance, lacking a direct
determination, we assumed the -enhancement as obtained by
Ca. The Eu line is so weak that it does not
require a hyperfine structure treatment.
The Li line at 6707 Å is an unresolved doublet, so we performed a spectrum synthesis analysis for this element as well.
For Eu and Li we took parameters for the 6707 Å line
from VALD and NIST
databases, and the
values
were simply averaged.
Chemical abundances obtained for the Sun are reported in Table 2 and compared with Grevesse & Sauval (1998). For Li we report the meteoritic value. We use our abundances instead of those from Grevesse & Sauval (1998) or Asplund et al. (2005) because they were obtained in an homogenous way with respect to our target stars.
Table 2:
Our measured solar abundances (
compared with Grevesse & Sauval
(1998).
Table 3:
Evolutionary stage and chemical abundances (
of our objects.
4.2 Continuum, atmospheric parameters and chemical abundance determination
The chemical abundances for Li, Ba, and Eu were obtained by comparing
observed spectra with synthetic ones (see Fig. 2 for an example of the
spectral synthesis applied to stars #236 and #591),
so in this case the continuum detrermination is a by-product
of this procedure.
For the remaining elements abundances were obtained from the equivalent
widths (EWs) of the spectral lines. In this second
case a continuum determination is a bit more complicated. Our spectra
are centered on the H region,
so they are relatively free from spectral lines, especially
compared to the data of F08, whose spectral range was much
bluer and contains many more lines.
This fact allowed us to proceed in the following way. First, we
selected for each line a region of 20 Å centered on the line
itself (this value is a good compromise between having enough
points, i.e. good statistics, and avoiding an excessively
large spectral region over which the spectrum could be substantially
curved). Then we built the histogram of the distribution of the flux
where the peak is a rough estimation of the continuum. We refined this
determination by fitting a
parabolic curve to the peak and using the vertex as our continuum
estimation. Finally, the continuum determination was revised by eye and
corrected by hand if a clear discrepancy with the spectrum was found.
Then, using the continuum value previously obtained, we fit
a Gaussian curve to each spectral line and obtained the EW
from integration. We rejected lines if affected by bad continuum
determination, by non-Gaussian shape, if their central wavelength did
not agree with that expected from our line-list, if the lines were too
broad or too narrow with respect to the mean FWHM, or if it was
affected by blending with other spectral features. We verified that the
Gaussian shape was a good approximation for our (mostly weak)
spectral lines, so generally no Lorentzian correction was
applied. The used lines and the measured EWs are reported
in Table 6.
Table 4: Internal errors associated to the chemical abundances.
Initial estimates of the atmospheric parameters were derived
from BVIJHK photometry.
Effective temperatures (
)
for each star were derived from
-[Fe/H]-color
relations (Ramírez
& Meléndez 2005; di Benedetto 1998; Alonso
et al. 1999). Colors were de-reddened assuming a
reddening value of 0.3 mag (Frinchaboy
et al. 2008). Then
was adjusted in order that abundances derived for individual
FeI lines showed no trend with excitation potential.
Surface gravities log g were obtained from
the canonical equation:

For


The luminosity
was obtained from the absolute magnitude MV
using the measured V magnitude, assuming
the bolometric correction (
)
from Alonso et al. (1999)
and an apparent distance modulus (m-M)V
obtained in the following way.
For many of the spectra we could measure the
FeII line at 6456 Å. While just one
FeII line is not
enough to determine a reliable gravity for a single star,
it can be used to determine the apparent distance modulus
(which is the same for all the stars), simply varying it until
the mean FeI and FeII abundances of the cluster (calculated
assuming the gravity obtained from the previous formula) are the same.
All other quantities in the gravity equation (
,
,
BC, V) are known.
We obtained:

We point out that this value was optimized in order to give [FeI/H] = [FeII/H], so it can differ from the real apparent distance modulus because the other variables in the gravity equation (mainly the mass) can be affected by systematic errors. This can explain the difference with F08, where the authors find


Finally, microturbulence velocity ()
was obtained from the relation:

which was obtained from Marino et al. (2008).
Adopted atmospheric parameters for each star are reported in Table 1. In this table Cols. 1 and 2 give the ID of the star according to Phelps et al. (1994, ID1) and our photometry (ID2). Columns 3 and 4 give the coordinates, Cols. 5-11 the B, V, I, J, H, K magnitudes, Col. 12 the temperature (K), Col. 13 the gravity, Col. 14 the microturbulence velocity (km s-1), and Col. 15 the heliocentric radial velocity (km s-1).
The Local Thermodynamic Equilibrium (LTE) program MOOG was used to determine the abundances, coupled with a model atmosphere interpolated from the Kurucz models for the parameters obtained as described before. Results are reported in Table 3.
4.3 Internal errors associated with the calculation of the chemical abundances
The abundances of every element are affected by measurement errors.
In this section our goal is to estimate the total amount of
the internal error (
)
affecting our data. Clearly, this requires an accurate analysis of all
the internal sources of errors. External errors can be estimated by
comparison with other works, as done in Sect. 6.
It must be noted that two sources of error mainly contribute
to
.
They are:
- the errors
due to the uncertainties in the EWs measures or in the comparison of observed spectra with synthetic ones;
- the uncertainty
introduced by errors in the atmospheric parameters adopted to compute the chemical abundances.







![]() |
Figure 3: Abundance ratios obtained for our stars. Typical error bars and mean values (dashed lines) are indicated. Open circles are values rejected during the sigma clipping rejection procedure (see Sect. 5). |
Open with DEXTER |
Errors in temperature were determined as in Marino
et al. (2008). The mean error
turned out to be
50 K.
The uncertainty in the surface gravity has been obtained by the
canonical formula using the propagation of errors. The variables used
in this formula that are affected by random errors are
and the V magnitude. For the temperature
we used the error previously obtained, while for V
we assumed a mean error of 0.1 mag, which is the
typical random error for stars at that magnitude. Other error sources
(distance modulus, reddening, bolometric correction, mass) affect
gravity in a systematic way, so are not important to our
analysis. The mean gravity error turned out to be 0.06 dex.
This implies a mean error
in the microturbulence of 0.02 km s-1.
Once the internal errors associated with the atmospheric parameters were calculated, we re-derived the abundances of one reference star (#591), assumed to represent our sample, by using the following combination of atmospheric parameters:
- (
,
,
);
- (
,
,
);
- (
,
,
),



The difference of abundance between values obtained with the
original and modified values gives
the errors in the chemical abundances due to uncertainties in each
atmospheric parameter. They are listed in Table 4 (Cols. 3-5).
Our best estimate of the total error associated to the abundance
measures is calculated as

listed in the Col. 6 of Table 4. Column 7 of Table 4 is the observed dispersion obtained as described in Sect. 5. Comparing


5 Results of abundance analysis and revised cluster parameters
The chemical abundances we obtained are summarized in Table 3 together with the
evolutionary stage of each star. This latter is based on the star
position in the CMD (Fig. 5),
where there is a clear division between RGB and RC stars. The
metallicity of the cluster turns out to be sub-solar with:
![\begin{eqnarray*}{\rm [Fe/H] =-0.31\pm0.02~dex}.
\end{eqnarray*}](/articles/aa/full_html/2010/01/aa13258-09/img74.png)
The abundance ratio trends versus [Fe/H] with the exception of Li are shown in Fig. 3. The bottom right panel report also the [Eu/Ba] content, which is an indicator of the r-process/s-process element ratio, because Eu is an almost pure r-element, while Ba is mainly produced in s-processes. Mean abundance ratios were calculated using the sigma clipping rejection method. Rejected measures are plotted as open circles. Mean values are reported in Fig. 3 and Table 3. The cluster turns out to have a solar scaled composition for Sc and Ni, while Ba is slightly overabundant and Eu is overabundant. Ca is overabundant, and its value (+0.16) is typical of a cluster of this metallicity. The Ca overabundance found here is consistent with the Ti overabundance found in F08, as expected from these


The mean [Eu/Ba] value turns out to be +0.17 dex.
This is an indication that Tombaugh 2 could be a intermediate
object between the thin and thick disk, according to Mashonkina et al. (2003, see
their Fig. 4).
In fact in that paper thick disk stars have
[Eu/Ba] greater than 0.35 dex, while thin disk stars
have [Eu/Ba] 0.0
on average. However, its large distance from the Galactic plane (
700 pc)
makes it a certain thick disk member. According to the measured
[Eu/Ba] value, stars of Tombaugh 2 were formed from
material more enriched by r-process elements with respect to the Sun,
being the cluster
1.5 times
richer in Eu that our star.
![]() |
Figure 4: Li abundances for observed stars. The evolutionary stage of each star is indicated. The symbols with arrows represent upper limits. Typical error bars are indicated. |
Open with DEXTER |
Li is a very fragile element, which is easily destroyed in stellar
interiors at relatively low temperature (2.5 106 K).
During the life of a star and in particular during the
MS phases, the Li-rich material that lies near the surface is
circulated downwards where the temperature is high enough for
Li burning to occur. When the star evolves to the red giant
phase, the deepening of the convective envelope brings up to the
surface internal matter which was nuclearly processed and the
Li abundance decreases. Several studies have shown that
further Li depletion occurs after the first dredge-up,
evidencing the action of an extra-mixing process (e.g. Charbonnel 1995). This can be seen
in Fig. 4,
where we plot log
(Li) vs.
temperature. RGB stars are indicated as filled circles, while
RC stars are open squares. One immediately notes that
RGB objects have a mean Li abundance greater than
RC stars, for which we could measure only upper limits
(see i.e. stars #236 and #591 in
Fig. 2).
Only the star #2846 falls out of this picture because it is
appears to be a RC object, but its
Li content is high (as high as the other
RGB stars). This object could be explained by error in the
photometry, which altered its real position in the CMD.
Otherwise, we found a clear trend of Li as a function of
for RGB stars. Apparently objects hotter than 4900 K
(and located in the lower RGB) have a constant Li
abundance (log
(Li)
1.18).
For stars colder than this limit Li starts to be more and more
depleted as the stars climb up the RGB (see dashed line in
Fig. 4).
A very similar behavior was found by De Freitas
et al. (in prep.) based on UVES spectra of
several open cluster giant members.
Having new values for metallicity and -element content, we
redetermined the basic parameters for the cluster using Padova
isochrones (Marigo et al. 2008)
and our V, I photometry.
The apparent distance modulus we found ((m-M)V = 15.1)
agrees well with the one from the FeI/II ionization
equilibrium. According to our fit (see Fig. 5), Tombaugh 2
has a reddening of E(B-V) =
0.25 (
E(V-I)=0.32)
and an age of 2.0
0.1 Gyr, which agrees well with the value given
by F08.
6 Comparison with previous metallicity studies
We first compare our findings with those of Brown et al. (1996) and Friel et al. (2002). The
former conducted a high resolution (R 34 000)
study of To2 with the CTIO 4-m telescope and found
[Fe/H] = -0.40
0.25 for E(B-V)=0.4
or [Fe/H]=-0.5
0.25 for E(B-V)=0.3
based on three stars, with [Fe/H] = -0.2, -0.4,
and -0.6. The Brown abundance analysis showed that To2 has a
reddening of E(B-V)=0.3-0.4.
The metallicities of their three stars are on the low end of the range
of our derived [Fe/H] values. However, the more metal-rich of
the three Brown et al. stars is ruled out as a cluster member
based on our radial velocity criterion. Thus, accounting for this, the
Brown et al. metallicity would be lower than our mean value,
by -0.2 to -0.3 dex, but still within 1
of our value.
Friel et al. (2002)
also obtained spectroscopic abundances for To2, using the CTIO
4-m/ARGUS, which yielded a much lower resolution (
)
than either our study, F08 or that of Brown
et al. (1996), and had metallicities determined from
spectral indices. Friel found [Fe/H] = -0.44
0.09 for To2 from a sample of 12 member stars and the
individual measurements ranging from-0.28 to -0.65. Their value agrees
well with ours.
![]() |
Figure 5: Isochrone fitting to the CMD. The derived parameters are indicated. |
Open with DEXTER |
We now turn to a comparison with F08. Recall that they also used the
VLT, obtaining both high resolution UVES data as well as lower
resolution GIRAFFE data. However, their UVES spectra were of
very low S/N, 15-20, while
GIRAFFE data were taken in a much more crowded region of the
spectrum - 4750-6800 Å - than we employ
here - 6500-6800 Å. Their GIRAFFE resolution (R=26 000)
is only slightly higher than ours. They derived detailed abundances for
18 cluster members (4 from UVES data),
i.e. with photometry consistent with being RGB or
RC stars and with
velocities in the range 121
4 km s-1. They found evidence
for two populations of To2 stars, with a metal-poor group
(mean [Fe/H] = -0.28) and a metal-rich group
(mean [Fe/H] = -0.06), with a fairly small dispersion
among each group. We refer to these as the MP and MR groups
subsequently. The total metallicity range covered a wide
margin: 0.00 to -0.43. Their data also suggested
distinct Ti abundances for these 2 groups, with the
MP group being
-enhanced
([Ti/Fe] = +0.36) and the MR group solar
([Ti/Fe] = +0.02). In addition, they found that the
MP group was more centrally concentrated than their
MR counterparts. They argued that their errors were unable to
account for these very surprising results and suggested that they were
indeed real. After discussing a number of possible scenarios to explain
their findings, F08 argued that the most likely was that the
centrally concentrated MP group represented the true
To2 cluster stars and that the MR group was a
spatially overlapping and kinematically associated but cold stellar
stream. They found this scenario to be more feasible than the even more
dubious possibility that To2 possessed a real metallicity
spread, which would make it unique among Milky Way open
clusters, but also claimed this as a viable alternative. Finally, they
associated To2 as a likely member of the
GASS/Mon stream.
Table 5: The 22 non-member stars and the two members with bad spectra.
Our results yield a mean metallicity of -0.31
0.02 from 13 members, with values varying from -0.19
to -0.39, all lying in the range of F08's MP group.
This variation is less than half of that covered by all of
F08 stars, already a strong hint that F08 may have
underestimated their errors. Of equal importance, we find no
hint of bimodality in the metallicity distribution. Our errors, which
we estimate to be of the same order as those estimated by F08,
can completely account for our observed Fe abundance
variation. What about stars in common between the two studies? There
are six, and the detailed comparison is as follows, where we give
(our ID: F08 ID, our [Fe/H]: F08 [Fe/H]).
(1672:135, -0.38:-0.07), (1827:140, -0.25:0.00), (238:127,
-0.30:-0.34), (3574:199, -0.23:-0.20), (3763:182, -0.38:-0.03), and
(591:164, -0.37:-0.11). Of these, all but the latter were both
observed with GIRAFFE. Only two of the stars show good agreement, while
the other four consistently have lower metallicities in our study than
what was found by F08, in particular the lone UVES star. All
four of these stars lie in their MR group, but now we find
that
actually they have metallicities that would have placed them in their
MP group. The differences are substantially larger than
expected, given the combined error estimates. Note that all four
UVES stars, all with very low S/N,
lie in the MR group.
Table 6: The linelist.
How do we explain the discrepancy between our results and
those of F08? Is it possible to explain it through
differences in the stellar parameters of their two groups? We indeed
find that their MP stars are on average about
50 K cooler, have a
some 0.3 smaller and a
0.05
less than their MR stars. The combined effects of these
differences on the derived metal abundance, based on their estimates of
this value to the input atmospheric parameters, can only explain about
0.05 dex of the 0.22 dex difference between the two
mean metallicities and is thus insufficient. We feel that the most
likely explanation is that F08 seriously underestimated their errors.
In particular, their UVES spectra are all of a very
low quality. On the other hand the spectral regions of
GIRAFFE spectra lie substantially blueward of our spectra,
in much more crowded regions. All this makes EWs more
difficult to be measured and atmospheric parameters consequently less
accurate than those obtained in this paper.
Thus, we do not agree with the main result of F08,
viz. that a real abundance spread exists among
To2 stars and/or stars associated with the purported stellar
stream. There is now no need for such a stream - all stars are
simply To2 members, and all have the same metallicity. Our
stars cover an even larger area than those of F08
(see Fig. 1),
yet there is no sign of any radial metallicity dependence,
another result they suggested. It appears that
F08 results on their MP group are actually applicable
to all To2 members. A further advantage of our
metallicity is that it agrees much better with that expected for a
cluster at the Galactocentric distance of To2, which in fact
was used as one of the arguments by F08 for preferring the
MP group as a representative of the cluster. As they
point out, stars as metal rich as those in their MR group
(nearly solar metallicity) at
15 kpc are not consistent with the measured metallicities
of red giants in the outer disk (Carney
et al. 2005) as well as outer disk Cepheid stars (Yong et al. 2006), which
both suggest that the median disk metallicity at this Galactocentric
distance should be [Fe/H]
-0.4, similar to our To2 value. These studies of outer disk
tracers find no stars as metal rich as the
F08 MR group, even for the younger Cepheid
populations. These outer disk stars also show enhanced
-element
abundances. Indeed, if F08 results for their MP stars
are correct, then To2 stars are Ti-enhanced (
[Ti/Fe]
= +0.36),
and we also find them to be Ca enhanced (
[Ca/Fe]
= +0.16), as
found by Carney et al. (2005)
and Yong et al. (2006)
for outer disk stars. However, we do not find any trend of
[Ca/Fe] vs. [Fe/H], as F08 found for Ti.
Again, evidence for a single chemical composition for
To2 stars, as we find, appears stronger and much more
tractable than any of the F08 alternatives. Finally, we note
that the slight Ba overabundance measured by us is in
perfect agreement with the results of D'Orazi
et al. (2009) who reported the discovery of a trend
of increasing barium abundance with decreasing age based on a large
sample of Galactic open clusters. Actually, the inclusion
of To2 in the sample allows for better sampling of the age
interval between 2 and 4 Gyr.
To conclude, we have no additional evidence for or against F08's contention that To2 is indeed a likely member of the GASS/Mon stream. However, their argument which read that added weight is given to this possibility by their observed To2 metallicity spread, is now rebuked. They argued that, since NGC 2808 is also believed to be a member and it shows evidence for an internal population dispersion (Piotto et al. 2007), the fact that To2 also shows such a dispersion favors its GASS/Mon membership.
7 Discussion and summary
The existence of multiple populations in Galactic globular clusters and
some large magellanic cloud clusters is now well established (Milone
et al. 2009; Piotto 2008; Marino et al. 2009,2008).
Indeed, Centauri
is now known to possess at least five populations showing different
ages and abundances (Villanova
et al. 2007). However, certainly not all such
clusters show this phenomenon, at least not in a clear way. Indeed, the
dominant characteristic that seems to be in common among such clusters
is their mass - only the most massive clusters are involved.
Given the current paradigm concerning the origin of these multiple
populations - retention of gas left over and polluted by a
first generation of star formation to form a second, chemically
distinct population - obviously requires sufficient cluster
mass to retain the required SN and AGB or massive star wind
ejecta. The required mass is estimated to be some 10
(Mieske et al. 2008).
Thus, only the most massive Galactic globulars fulfill this requirement
(although note that at least one less massive cluster, M4 at 10
,
also shows multiple populations, Marino
et al. 2008). The case of the LMC clusters
is more uncertain, but it is clear that the phenomenon is
correlated with mass. On the other hand, Galactic open clusters
generally do not exceed masses of
,
so that the possibility of retaining any first generation ejecta is
very unlikely - any gas that remains after the first star
formation episode is subsequently quickly removed by the ejecta itself.
Thus, one does not expect to find multiple populations in open
clusters. To2 is a typical open cluster in this respect and
certainly does not exceed a few
.
The finding by F08 that To2 possessed an intrinsic metallicity spread (or was a superposition of an outer disk cluster well out of the plane and a cold stellar stream with exactly the same velocity but distinct chemistry) was then understandably met with some incredulity and at the very least required independent corroboration. Hence the present study.
We obtained independent data, but used the same telescope and spectrograph as they did. However, we deliberately achieved superior data quality - in both S/N and by observing in a longer wavelength regime, where line crowding was significantly reduced. Our reduction and analysis procedures were virtually identical to theirs.
To summarize, we obtained high resolution VLT+GIRAFFE spectra in the optical for a number of stars in the Tombaugh 2 field. Radial velocities and the position in the color-magnitude diagram are used to assess cluster membership. The spectra, together with input atmospheric parameters and model atmospheres, are used to determine detailed chemical abundances for a variety of elements in 13 stars confirmed as members.
We derive a mean metallicity [Fe/H] = -0.31
0.02, with no evidence for an intrinsic abundance dispersion. We find
Ca, Ba, and Eu to be enhanced while Ni and Sc are solar.
Li abundances decrease with
on the upper giant branch and maintain a low level for red clump stars.
The mean metallicity we derive agrees well with that expected from the
radial abundance gradient in the disk for a cluster at its
Galactocentric distance. The surprising result found by F08,
viz. two distinct abundance groups within the cluster,
implying either a completely unique open cluster
with an intrinsic metallicity spread, or a very unlikely superposition
of a cold stellar stream and a very distant open cluster,
is not supported by our new data. We suspect that the
F08 data was subject to substantially larger errors than they
estimated, especially given the low S/N
of their UVES spectra and their much bluer wavelength range,
which was plagued by line crowding. To2, instead of being a
unique cluster, is found to be a normal representative of
its class.
E.C., S.V. and D.G. gratefully acknowledge support from the Chilean Centro de Astrofísica FONDAP No. 15010003. E.C. and D.G. arealso supported by the Chilean Centro de Excelencia en Astrofísica y Tecnologías Afines (CATA).
References
- Alonso, A., Arribas, S., & Martínez-Roger, C. 1999, A&AS, 140, 261 [Google Scholar]
- Asplund, M., Grevesse, N., & Sauval, A. J. 2005, in Cosmic Abundances as Records of Stellar Evolution and Nucleosynthesis, ed. T. G. Barnes III, & F. N. Bash, ASP Conf. Ser., 336, 25 [Google Scholar]
- Bellazzini, M., Ibata, R., Monaco, L., et al. 2004, MNRAS, 354, 1263 [NASA ADS] [CrossRef] [Google Scholar]
- Brown, J. A., Wallerstein, G., Geisler, D., & Oke, J. B. 1996, AJ, 112, 1551 [NASA ADS] [CrossRef] [Google Scholar]
- Butler, D. J., Martínez-Delgado, D., Rix, H.-W., Peñarrubia, J., & de Jong, J. T. A. 2007, AJ, 133, 2274 [NASA ADS] [CrossRef] [Google Scholar]
- Carney, B. W., Yong, D., Teixera de Almeida, M. L., & Seitzer, P. 2005, AJ, 130, 1111 [NASA ADS] [CrossRef] [Google Scholar]
- Carraro, G., & Costa, E. 2009, A&A, 493, 71 [Google Scholar]
- Carraro, G., Bresolin, F., Villanova, S., et al. 2004, AJ, 128, 1676 [NASA ADS] [CrossRef] [Google Scholar]
- Carraro, G., Vázquez, R. A., Moitinho, A., & Baume, G. 2005, ApJ, 630, L153 [NASA ADS] [CrossRef] [Google Scholar]
- Carraro, G., Geisler, D., Villanova, S., Frinchaboy, P. M., & Majewski, S. R. 2007, A&A, 476, 217 [Google Scholar]
- Charbonnel, C. 1995, ApJ, 453, L41 [NASA ADS] [CrossRef] [Google Scholar]
- Chiappini, C., Matteucci, F., & Romano, D. 2001, ApJ, 554, 1044 [NASA ADS] [CrossRef] [Google Scholar]
- Crane, J. D., Majewski, S. R., Rocha-Pinto, H. J., et al. 2003, ApJ, 594, L119 [NASA ADS] [CrossRef] [Google Scholar]
- di Benedetto, G. P. 1998, A&A, 339, 858 [Google Scholar]
- Dinescu, D. I., Martínez-Delgado, D., Girard, T. M., et al. 2005, ApJ, 631, L49 [NASA ADS] [CrossRef] [Google Scholar]
- D'Orazi, V., Magrini, L., Randich, S., et al. 2009, ApJ, 693, L31 [NASA ADS] [CrossRef] [Google Scholar]
- Friel, E. D., Janes, K. A., Tavarez, M., et al. 2002, AJ, 124, 2693 [NASA ADS] [CrossRef] [Google Scholar]
- Frinchaboy, P. M., Majewski, S. R., Crane, J. D., et al. 2004, ApJ, 602, L21 [NASA ADS] [CrossRef] [Google Scholar]
- Frinchaboy, P. M., Marino, A. F., Villanova, S., et al. 2008, MNRAS, 391, 39 [Google Scholar]
- Grevesse, N., & Sauval, A. J. 1998, Space Sci. Rev., 85, 161 [NASA ADS] [CrossRef] [Google Scholar]
- Ibata, R. A., Irwin, M. J., Lewis, G. F., Ferguson, A. M. N., & Tanvir, N. 2003, MNRAS, 340, L21 [NASA ADS] [CrossRef] [Google Scholar]
- Kupka, F., Piskunov, N., Ryabchikova, T. A., Stempels, H. C., & Weiss, W. W. 1999, A&AS, 138, 119 [Google Scholar]
- Kurucz, R. L. 1992, in The Stellar Populations of Galaxies, ed. B. Barbuy, & A. Renzini, IAU Symp., 149, 225 [Google Scholar]
- Magrini, L., Sestito, P., Randich, S., & Galli, D. 2009, A&A, 494, 95 [Google Scholar]
- Marigo, P., Girardi, L., Bressan, A., et al. 2008, A&A, 482, 883 [Google Scholar]
- Marino, A. F., Villanova, S., Piotto, G., et al. 2008, A&A, 490, 625 [Google Scholar]
- Marino, A. F., Milone, A. P., Piotto, G., et al. 2009, A&A, 505, 1099 [Google Scholar]
- Martin, N. F., Ibata, R. A., Bellazzini, M., et al. 2004, MNRAS, 348, 12 [NASA ADS] [CrossRef] [Google Scholar]
- Martínez-Delgado, D., Butler, D. J., Rix, H.-W., et al. 2005, ApJ, 633, 205 [NASA ADS] [CrossRef] [Google Scholar]
- Mashonkina, L., Gehren, T., Travaglio, C., & Borkova, T. 2003, A&A, 397, 275 [Google Scholar]
- Mieske, S., Hilker, M., Jordán, A., et al. 2008, A&A, 487, 921 [Google Scholar]
- Milone, A. P., Bedin, L. R., Piotto, G., & Anderson, J. 2009, A&A, 497, 755 [Google Scholar]
- Moitinho, A., Vázquez, R. A., Carraro, G., et al. 2006, MNRAS, 368, L77 [NASA ADS] [CrossRef] [Google Scholar]
- Momany, Y., Zaggia, S. R., Bonifacio, P., et al. 2004, A&A, 421, L29 [Google Scholar]
- Momany, Y., Zaggia, S., Gilmore, G., et al. 2006, A&A, 451, 515 [Google Scholar]
- Newberg, H. J., Yanny, B., Rockosi, C., et al. 2002, ApJ, 569, 245 [NASA ADS] [CrossRef] [Google Scholar]
- Pasquini, L., Avila, G., Blecha, A., et al. 2002, The Messenger, 110, 1 [NASA ADS] [Google Scholar]
- Peñarrubia, J., Martínez-Delgado, D., Rix, H. W., et al. 2005, ApJ, 626, 128 [NASA ADS] [CrossRef] [Google Scholar]
- Phelps, R. L., Janes, K. A., & Montgomery, K. A. 1994, AJ, 107, 1079 [NASA ADS] [CrossRef] [Google Scholar]
- Piotto, G. 2008, Mem. Soc. Astron. Ital., 79, 334 [Google Scholar]
- Ramírez, I., & Meléndez, J. 2005, ApJ, 626, 465 [Google Scholar]
- Randich, S., Bragaglia, A., Pastori, L., et al. 2005, The Messenger, 121, 18 [Google Scholar]
- Rocha-Pinto, H. J., Majewski, S. R., Skrutskie, M. F., & Crane, J. D. 2003, ApJ, 594, L115 [NASA ADS] [CrossRef] [Google Scholar]
- Skrutskie, M. F., Cutri, R. M., Stiening, R., et al. 2006, AJ, 131, 1163 [NASA ADS] [CrossRef] [Google Scholar]
- Twarog, B. A., Ashman, K. M., & Anthony-Twarog, B. J. 1997, AJ, 114, 2556 [NASA ADS] [CrossRef] [Google Scholar]
- Villanova, S., Piotto, G., King, I. R., et al. 2007, ApJ, 663, 296 [NASA ADS] [CrossRef] [Google Scholar]
- Villanova, S., Piotto, G., & Gratton, R. G. 2009, A&A, 499, 755 [Google Scholar]
- Yanny, B., Newberg, H. J., Grebel, E. K., et al. 2003, ApJ, 588, 824 [NASA ADS] [CrossRef] [Google Scholar]
- Yong, D., Carney, B. W., Teixera de Almeida, M. L., & Pohl, B. L. 2006, AJ, 131, 2256 [NASA ADS] [CrossRef] [Google Scholar]
Footnotes
- ... 2
- Based on observations collected at ESO-VLT, Paranal Observatory, Chile, Program numbers 076.D-0220(A).
- ...
SPECTRUM
- See http://www.phys.appstate.edu/spectrum/spectrum.html for more details.
- ... MOOG
- See http://verdi.as.utexas.edu/moog.html for more details
- ... NIST
- NIST database can be found at http://physics.nist.gov/PhysRefData/ASD/lines_form.html
All Tables
Table 1: Target stars and parameters.
Table 2:
Our measured solar abundances (
compared with Grevesse & Sauval
(1998).
Table 3:
Evolutionary stage and chemical abundances (
of our objects.
Table 4: Internal errors associated to the chemical abundances.
Table 5: The 22 non-member stars and the two members with bad spectra.
Table 6: The linelist.
All Figures
![]() |
Figure 1: Left panel: the CMD of Tombaugh 2 in the optical and infrared. Members stars studied in this paper are indicated as filled circles. Right panel: distribution on the sky of our targets (filled points) and F08 objects (empty squares). |
Open with DEXTER | |
In the text |
![]() |
Figure 2:
Example of the spectral synthesis method applied to stars #236
and #591. Observed spectra are plotted with continuous lines,
while the synthetic ones are plotted with dashed lines. Abundances (
|
Open with DEXTER | |
In the text |
![]() |
Figure 3: Abundance ratios obtained for our stars. Typical error bars and mean values (dashed lines) are indicated. Open circles are values rejected during the sigma clipping rejection procedure (see Sect. 5). |
Open with DEXTER | |
In the text |
![]() |
Figure 4: Li abundances for observed stars. The evolutionary stage of each star is indicated. The symbols with arrows represent upper limits. Typical error bars are indicated. |
Open with DEXTER | |
In the text |
![]() |
Figure 5: Isochrone fitting to the CMD. The derived parameters are indicated. |
Open with DEXTER | |
In the text |
Copyright ESO 2010
Current usage metrics show cumulative count of Article Views (full-text article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.
Data correspond to usage on the plateform after 2015. The current usage metrics is available 48-96 hours after online publication and is updated daily on week days.
Initial download of the metrics may take a while.