A&A 407, 289-301 (2003)
P. Sestito1 - S. Randich2 - J.-C. Mermilliod3 - R. Pallavicini4
1 - Dipartimento di Astronomia, Università di Firenze, Largo E. Fermi 5, 50125 Firenze, Italy
2 - INAF/Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy
3 - Institut d' Astronomie, Université de Lausanne, 1290, Chavannes-des-Bois, Switzerland
4 - INAF/Osservatorio Astronomico di Palermo, Piazza del Parlamento 1, 90134 Palermo, Italy
Received 27 February 2003 / Accepted 12 May 2003
We have carried out a high resolution spectroscopic survey of the 220-250 Myr old cluster NGC 6475: our main purpose is to investigate Li evolution during the early stages of the Main Sequence. We have determined Li abundances for 33 late F to K-type X-ray selected cluster candidates, extending the samples already available in the literature; for part of the stars we obtained radial and rotational velocities, allowing us to confirm the membership and to check for binarity. We also estimated the cluster metallicity which turned out to be over-solar ( ). Our Li analysis evidenced that (i) late F-type stars ( 6000 K) undergo a very small amount of Li depletion during the early phases on the ZAMS; (ii) G-type stars (6000 5500 K) instead do deplete lithium soon after arrival on the ZAMS. Whereas this result is not new, we show that the time scale for Li depletion in these stars is almost constant between 100 and 600 Myr; (iii) we confirm that the spread observed in early K-type stars in younger clusters has converged by 220 Myr. No constraints can be put on later-type stars. (iv) Finally, we investigate the effect of metallicity on Li depletion by comparing NGC 6475 with the similar age cluster M 34, but we show that the issue remains open, given the uncertain metallicity of the latter cluster. By using the combined NGC 6475+M 34 sample together with the Hyades and the Pleiades, we compare quantitatively Li evolution from the ZAMS to 600 Myr with theoretical predictions of standard models.
Key words: stars: abundances - stars: evolution - open clusters and associations: individual: NGC 6475
Standard models predict a certain amount of depletion during Pre-Main Sequence (PMS) evolution of solar-type stars (e.g. D' Antona & Mazzitelli 1994) and no depletion at all after arrival on the Zero-Age Main Sequence (ZAMS). The predicted PMS Li depletion increases as mass decreases while the predicted Main Sequence (MS) depletion remains small in all cases, apart from the coolest stars; moreover Li depletion should depend only on age, chemical composition and mass (or effective temperature), i.e. stars with the same mass in a given cluster should all have the same Li abundance. As far as the Li-metallicity relationship is concerned, standard models predict that increased metal abundances should lead to a significant increase of Li depletion during PMS contraction for stars cooler than 6000 K (e.g. Chaboyer et al. 1995; Swenson et al. 1994): this is due to the fact that the gas opacity grows up in stars with a higher iron content, thus the depth of the convective zone (CZ) increases leading to a large amount of Li depletion. It is worth of mention that also oxygen, as well as other elements, largely contributes to affect the opacity values and thus the depth of the CZ (Piau & Turck-Chièze 2002).
The predictions of standard models are in contrast with observational results. Namely, focusing on the evolution of Li up to the Hyades age (600 Myr), observations of very young clusters (ages 30-50 Myr) show that solar analogs undergo very little (if any) Li depletion during the PMS (e.g. Martín & Montes 1997; Randich et al. 1997); on the other hand, the comparison of clusters of different ages clearly shows that these stars do deplete Li while on the MS. Furthermore, the star-to-star scatter in Li abundance seen in young clusters for stars cooler than 5500 K (e.g. Soderblom et al. 1993 - hereafter S93; García López et al. 1994; Jones et al. 1996; Randich et al. 1998) is clearly in contrast with standard model predictions; this dispersion is already present at arrival on the ZAMS and has disappeared by the age of the Hyades (600-700 Myr, Thorburn et al. 1993). Based on the observed Li-rotation relationship (S93), and specifically on the fact that rapid rotators in young clusters have on average higher Li abundances than slow rotators, the most commonly accepted explanation for the scatter is that Li depletion is connected to rotationally driven mixing and angular momentum transport. Note however that slow rotators in young clusters can have either high or low Li abundances, as showed by Randich et al. (1998) for Persei, and that the Li-rotation relationship breaks down for the coolest stars ( 4500 K, García López et al. 1994).
Finally, the fact that no Li depletion-metallicity relationship has so far been convincingly demonstrated (see for example Jeffries & James 1999), at least for 4700 K, appears in contrast with model predictions.
In summary, the question remains which mechanism(s) drives or inhibits Li depletion in stars of different masses during the PMS and MS phases.
In order to put additional empirical constraints on early-MS Li depletion processes, it is necessary to enlarge the database of Li observations in young clusters; for this reason, we carried out a survey of NGC 6475, a well populated Southern hemisphere cluster, with reported over-solar metallicity. NGC 6475 is a very good target, since its age of 220 Myr (Meynet et al. 1993) is intermediate between those of the Pleiades and of the Hyades and it is the closest and most compact open cluster at that age (distance 250 pc); the estimated spectroscopic iron abundance is and the reddening E(B-V)=0.06 (James & Jeffries 1997 - hereafter JJ97). Li data for this cluster allow us to investigate early-MS Li depletion and its time scale for solar-type and lower mass stars, as well as, to some extent, the dependence of Li depletion on metallicity by comparing NGC 6475 to M 34 (NGC 1039), surveyed by Jones et al. (1997). The latter cluster is about co-eval to NGC 6475 (250 Myr, Jones & Prosser 1996); Canterna et al. (1979) found a solar metallicity for M 34, based on multicolor ubvy photometry of two F-type stars, while a recent and more detailed high resolution spectroscopic analysis by Schuler et al. (2003) evidenced that the iron content could be over-solar ( ). Note that the result of Schuler et al. (2003) is based on five solar-type stars; had they considered their whole sample of nine stars with K, they would have found (see the quoted reference for more details).
Previous studies of Li in NGC 6475 were carried out by JJ97 and James et al. (2000) (hereafter J00): the first one is based on an X-ray selected sample, while in the latter one an optically selected sample is studied. In this paper, we present the results of additional Li observations of NGC 6475: our data, merged with the ones of the two previous works, provide a larger and more statistically significant sample of stars to further address the issue of Li evolution between the age of Pleiades and that of the Hyades. In addition, our sample contains a few more stars of later spectral-types than the previous surveys, allowing us to get some insights on early-MS Li depletion for the coolest stars.
In Sect. 2 we describe our sample and the observations; in Sect. 3 we summarize the radial velocities and abundances analysis, while the results and a discussion are presented in Sects. 4 and 5. Finally our conclusions (Sect. 6) close the paper.
Our original sample includes 34 cluster candidates with , selected from the X-ray survey of Prosser et al. (1995) (hereafter P95). Target stars and photometry are listed in the first three columns of Table 1: the identification number (Col. 1) and the photometry (Cols. 2 and 3) were retrieved from P95; a reddening E(B-V)=0.06 was adopted.
The observations were carried out during three runs (April 1994, April 1995 and July 1996) at the European Southern Observatory (ESO), La Silla, Chile, with the 3.6 m telescope equipped with CASPEC. During the April 1994 and April 1995 observing runs the standard echelle grating (31.6 lines ) with the red cross-disperser (158 lines ) and the short camera were used, together with ESO CCD #32 (TK512, with pixels of 27 ); the nominal resolving power was (slit aperture of 280 ). In the July 1996 run the long camera and the ESO CCD #37 (TK1024, with pixels of 24 ) were used with a slit aperture of 200 , which provided a slightly larger resolving power, . Exposure times ranged between 10 min and 1.5 hours resulting in S/N ratios of 50-150. For each star, the corresponding CASPEC run and exposure time are listed in Cols. 4 and 5 of Table 1.
About 80% of the stars were also observed with the CORAVEL instrument (Baranne et al. 1979) at the 1.54 m Danish telescope at ESO. Three observations were obtained through the period 1985-1995 for stars belonging to a program designed to search for faint members, and one or two measurements were obtained in April and July 1996 for stars selected from the ROSAT X-ray source catalog (P95) during the course of a long-term systematic program on cluster red dwarfs started in Chile in 1983. We mention that part of the stars of the JJ97 sample were also observed by us with CORAVEL. The CORAVEL observations allowed us to derive radial velocities and to check for binarity; for part of the stars the projected rotational velocities () were also determined.
CASPEC data reduction was performed with the package MIDAS in the ECHELLE context,
following the usual steps: bias subtraction, order definition, order extraction, inter-order background subtraction, flat-fielding and wavelength calibration.
Figure 1 shows examples of normalized spectra
in the Li region.
|No.||V||(B-V)0||Observing run||Exposure time||RV||N||Spectroscopic membership|
|P95||(CASPEC)||[s]||[km s-1]||[km s-1]||(CORAVEL)|
|R10A||12.99||1.04||July 1996||3600||-||-||-||No information|
|R14||12.06||0.60||April 1995||1200||-15.60 0.24||4||6.0 1.6||M|
|R15A||12.28||0.65||July 1996||2700||+18.60 0.56||1||15.0 2.0||N?, SB|
|R16A||11.69||0.66||April 1994||720||-13.77 0.34||2||5.8 2.4||M|
|R26A||14.89||1.44||July 1996||5400||-||-||-||No information|
|R27||12.12||0.62||April 1994||1200||-14.98 0.33||3||9.8 1.3||M|
|R35||14.54||1.25||July 1996||4800||-||-||-||No information|
|R39A||12.36||0.65||April 1995||1320||-14.03 0.23||4||4.5 1.5||M|
|R42||11.83||0.57||July 1996||1800||-11.23 0.43||2||14.3 2.3||M, SB|
|R48A||14.73||1.22||July 1996||5400||-||-||-||No information|
|R49A||12.37||0.63||July 1996||2700||-13.90 0.50||1||5.9 2.8||M|
|R51A||12.25||0.63||April 1995||1200||-11.81 1.06||1||14.3 3.3||M|
|R53||13.40||0.84||April 1995||1800||-12.83 0.56||1||3.7 3.2||M|
|R55A||12.33||0.77||April 1995||1200||-19.58 0.47||1||9.0 1.8||M, SB?|
|R55B||11.70||0.54||July 1996||2400||-13.59 0.59||4||13.7 4.6||M, SB|
|R64||11.96||0.67||April 1994||1200||-14.44 0.56||2||5.8 1.8||M|
|R66||12.78||0.72||April 1995||1500||-15.71 0.45||3||7.7 1.3||M, SB2|
|R73A||10.95||0.51||April 1994||600||- 12.56 1.11||4||18.7 2.2||M, SB|
|R79A||11.92||0.70||July 1996||1800||-15.71 0.33||2||6.9 1.5||M|
|R82||12.87||0.67||April 1995||1620||-14.69 0.30||3||6.7 1.5||M|
|R92||12.46||0.63||April 1995||1020||+4.31 1.03||2||8.1 1.6||M, SB?|
|R97||12.17||0.63||April 1994||900||-13.97 0.50||1||9.7 1.8||M|
|R102||13.32||0.83||April 1995||2100||-15.97 0.56||1||6.1 4.3||M|
|R103||12.37||0.65||April 1995||1500||-13.56 0.43||1||2.4 2.8||M|
|R105||12.26||0.62||April 1994||1020||-10.22 0.53||1||11.0 2.0||M, SB|
|R109A||12.56||0.71||April 1994||1200||-6.31 1.06||2||6.8 3.0||M?, SB|
|R116||12.77||0.82||April 1995||1560||+10.59 25.86||3||7.5 1.8||M?, SB|
|R123||13.17||0.79||April 1995||1800||-15.37 0.54||1||9.8 2.1||M|
|R126A||11.45||0.50||April 1994||720||-||-||-||No information|
|R133||12.17||0.65||April 1995||1200||-12.72 16.07||3||7.5 1.3||M?, SB|
|R135A||13.08||0.96||July 1996||3300||-||-||-||No information|
|R136A||11.43||0.70||July 1996||1800||-14.83 0.81||1||26.2 2.6||M, SB2|
|R137A||13.43||0.93||July 1996||3600||-||-||-||No information|
|R140B||12.38||0.76||July 1996||3000||-14.83 0.44||1||1.9 3.1||M|
Results for further stars and individual observations will be discussed in a separate paper devoted to the study of NGC 6475 based on CORAVEL observations. If necessary, data used in the present paper are available from J.-C. Mermilliod.
CORAVEL radial velocities have been used in conjunction with available data,
from P95 and JJ97, which are usually based on one observation,
to detect spectroscopic binaries and confirm the membership of the other
stars. The results are recorded in Col. 9 of
Table 1 and Col. 5 of Table 2
as membership determinations and remarks on duplicity
(M: member, M?: possible member, N?: doubtful member,
N: non-member, SB: spectroscopic binary, SB?: possible
spectroscopic binary, No information: the
star has not been observed with CORAVEL). In Table 2
the identification numbers of JJ97 are used (Col. 1).
|JJ97||[km s-1]||[km s-1]||(CORAVEL)|
|1||-15.79 0.33||3||11.6 1.1||M|
|3||-15.57 2.36||2||65.8 13.4||M|
|6||-14.49 0.50||3||19.4 1.2||M|
|7||-19.15 0.45||2||12.0 1.5||M, SB|
|8||-14.74 0.54||1||4.3 2.9||M|
|19||-12.71 0.47||1||5.2 3.1||M|
|22||-15.37 0.53||3||17.3 1.5||M|
|24||-13.24 0.53||1||7.8 2.9||M|
|27||-11.19 0.59||1||15.5 2.1||M, SB2|
|29||-13.45 0.65||3||4.2 1.9||M, SB?|
|31||+34.44 7.37||2||19.2 1.3||M, SB|
|34||-6.31 1.96||2||6.8 3.0||M, SB|
|36||-24.67 3.24||2||12.6 2.4||M, SB|
Several stars require a comment:
We end this section with two comments: first, the high rate of confirmed members shows that X-ray surveys are effective in detecting new cluster members, not only for very young clusters like IC 2602 and IC 2391 (e.g. Randich et al. 2001), but also for somewhat older clusters; second, whereas a detailed discussion of the evolution of rotation is beyond the scope of this paper, we note that the projected rotational velocities are rather low for the majority of the stars. More specifically, considering our sample (Table 1) and the stars from the sample of JJ97 (Table 2), there are 26 stars (63%) with km s-1, 13 stars (32%) with km s-1and only 2 stars (5%) with km s-1: star R136A ( km s-1) and star JJ3 ( km s-1).
We assumed for all the sample stars the same surface gravity , while microturbulence was derived as (Nissen 1981, Boesgaard & Friel 1990); these two parameters have negligible effects on Li abundances, while they affect metallicity. The assumed random errors are 0.3 dex in and 0.3 km s-1 in .
We measured equivalent widths (EWs) of the Li I
doublet; at our
resolutions, this spectral feature is blended,
or partially blended, with the
line, therefore the
contribution of the latter feature needs to be considered.
The EW of the iron line was estimated
following the prescription of S93, namely
Note that the EW of star R26A
= 3860 K), marked with one asterisk in Col. 3 of Table 3,
has not been corrected
for the Fe I contribution, since the S93 formula is no more valid for
below 4000 K (see discussion
in Randich et al. 2000).
In this case, the quoted (Li) should be regarded as
an upper limit.
|P95||||[ ]||[ ]|
|R10A||4507||153 5||135 5||1.63 0.15|
|R14||5888||107 6||98 6||2.71 0.10|
|R16A||5656||155 10||145 10||2.77 0.17|
|R27||5810||122 3||113 3||2.73 0.10|
|R35||4095||73 18||51 18||0.28 0.21|
|R39A||5696||108 5||98 5||2.56 0.10|
|R42||6008||115 4||107 4||2.85 0.10|
|R48A||4144||41 10||20 10||-0.12 0.28|
|R49A||5772||92 6||82 6||2.53 0.10|
|R51A||5772||46 5||36 5||2.13 0.11|
|R53||5048||88 12||74 12||1.84 0.15|
|R55A||5272||86 10||74 10||2.05 0.14|
|R55B||6131||91 3||83 3||2.80 0.09|
|R64||5622||77 5||67 5||2.30 0.11|
|R66||5442||139 5||128 5||2.51 0.11|
|R73A||6257||63 15||56 15||2.69 0.17|
|R79A||5513||142 8||131 8||2.59 0.12|
|R82||5622||128 5||118 5||2.61 0.10|
|R92||5772||110 5||100 5||2.64 0.10|
|R97||5772||120 8||110 8||2.69 0.11|
|R102||5079||81 5||67 5||1.82 0.13|
|R103||5696||111 5||101 5||2.58 0.10|
|R105||5810||115 3||106 3||2.70 0.10|
|R109A||5477||121 7||110 7||2.43 0.12|
|R116||5110||83 4||70 4||1.87 0.12|
|R123||5206||107 6||94 6||2.12 0.12|
|R126A||6300||100 6||93 6||2.98 0.09|
|R133||5696||145 5||135 5||2.76 0.10|
|R135A||4706||170 3||154 3||1.76 0.14|
|R136A||5513||108 6||97 6||2.41 0.11|
|R137A||4787||156 3||140 3||1.79 0.13|
|R140B||5305||105 3||93 3||2.20 0.11|
|61||1||11.58||0.62||5810||124 4||2.79 0.10|
|42||2*||11.38||0.57||6008||105 9||2.84 0.11|
|81||3||11.38||0.48||6386||58 13||2.80 0.15|
|69||4||12.13||0.60||5888||30 10||2.14 0.33|
|127B||6||11.67||0.57||6008||90 14||2.75 0.13|
|127A||7||11.88||0.61||6330||101 9||3.05 0.10|
|94||8||13.30||0.81||5142||126 10||2.23 0.13|
|82||9*||12.87||0.67||5622||127 11||2.66 0.12|
|82B||10||12.53||0.68||5585||109 9||2.54 0.12|
|53||11*||13.40||0.84||5048||68 12||1.79 0.16|
|27||12*||12.12||0.62||5810||101 8||2.67 0.11|
|16A||13*||11.69||0.66||5659||102 7||2.56 0.11|
|7B||14||13.78||1.02||4555||123 17||1.62 0.17|
|7A||15||14.15||1.00||4604||203 25||2.04 0.21|
|39A||16*||12.36||0.65||5696||134 13||2.75 0.13|
|-||17||12.30||0.69||5549||91 11||2.40 0.13|
|-||18||11.99||0.71||5477||124 10||2.52 0.12|
|76||19||13.55||0.89||4899||104 11||1.88 0.14|
|-||20||14.05||0.98||4654||59 12||1.32 0.17|
|24||22||11.11||0.44||6564||60 8||2.90 0.15|
|103||23*||12.37||0.65||5696||99 12||2.57 0.13|
|33||24||13.42||0.83||5079||127 18||2.18 0.16|
|102||25*||13.32||0.83||5079||108 34||2.08 0.25|
|14||26*||12.06||0.60||5888||97 6||2.70 0.10|
|104||27||12.37||0.76||5305||126 8||2.38 0.12|
|72||28||10.79||0.49||6343||59 8||2.77 0.11|
|95||29||12.19||0.62||5810||97 8||2.64 0.11|
|39B||31||12.19||0.63||5772||118 11||2.73 0.12|
|119A||33||12.92||0.75||5339||103 9||2.29 0.12|
|109||34||12.63||0.74||5373||90 8||2.25 0.13|
|66||35*||12.78||0.72||5442||87 7||2.29 0.12|
|132||36||11.81||0.62||5810||49 14||2.29 0.17|
|1||40||13.30||1.07||4439||17 9||1.84 0.31|
|-||41||11.57||0.52||6214||89 8||2.89 0.10|
|134||42||12.41||0.78||5239||13 4||1.21 0.18|
Li abundances ((Li)) were computed from the measured EWs by interpolating
the curves of growth (COG) of S93; Li abundances were then corrected for
non local thermodynamic equilibrium (NLTE) effects by using the code
of Carlsson et al. (1994): NLTE corrections are provided only
= 4500 K. Whereas for cooler stars we adopted
LTE Li abundances, NLTE corrections are small
below 4500 K
(e.g. Pavlenko et al. 1995).
The measured EWs of the Li+Fe feature and the EWs corrected for the Fe I blend
are listed in Cols. 3 and 4 of Table 3, while
the derived Li abundances are listed in Col. 5.
Uncertainties in (Li) were computed by quadratically adding
the errors due to uncertainties in
and in EWs.
|103||12.55||0.66||5658||123 15||2.66 0.14|
|104||12.18||0.72||5442||192 13||2.80 0.14|
|105||12.92||0.77||5272||34 15||1.67 0.25|
|106||12.20||0.64||5734||148 11||2.84 0.12|
|109||12.07||0.67||5622||156 35||2.79 0.24|
|110||12.96||0.80||5174||90 13||2.06 0.15|
|111||11.59||0.52||6214||132 13||3.12 0.12|
|113||11.80||0.59||5928||133 17||2.92 0.14|
|115||11.70||0.60||5888||137 13||2.91 0.12|
|116||12.29||0.64||5734||158 12||2.88 0.12|
In the following, we will compare our results with those of JJ97 and J00 for NGC 6475 and with those of other clusters. In order to put all the data on a homogeneous scale, we recomputed Li abundances using the procedure described above and starting from published EWs for these clusters (NGC 6475 - JJ97 and J00; M 34 - Jones et al. 1997; Pleiades - S93 and Jones et al. 1996; and Hyades - Thorburn et al. 1993). Where necessary we also recomputed the effective temperatures (i.e. when the vs. B-V calibrations used by other authors were different from the used here).
Stars, photometry, EWs and (Li) for the samples of JJ97 and J00 are listed in Tables 4 and 5. The listed in in Tables 4 and 5 are those recomputed by us, after retrieving B-V values from P95. The listed EWs were corrected for the Fe I blend by JJ97 and J00, using a spectral subtraction technique. Ten of the stars of JJ97 (marked with asterisks in Table 4) are in common with our sample: for this reason, in the first two columns of Table 4, we report both the star numbers of P95 and of JJ97, while in Table 5 only the JJ97 numbering system is adopted.
All the stars in Table 4, with the exception of JJ40, JJ41 and JJ42, were considered as bona fide cluster members (see JJ97 and Table 2). The membership of JJ40, JJ41, JJ42 requires confirmation because the available velocities are based on a single observation and differ by more than 18 km s-1 from the cluster mean. If members, they clearly should be binaries.
Stars in Table 5 were selected by us following the same criteria adopted by J00, i.e. stars with radial velocities far from the cluster mean velocity were rejected; since star JJ105, rejected by J00, has a RV differing only by 5 km s-1 from the cluster mean velocity, we consider this object as a possible member.
|6703.57 Å||6725.36 Å||6726.67 Å||6750.16 Å||6810.27 Å||6820.37 Å||6828.60 Å||6858.15 Å||average( )|
|R14||-||7.87||7.61||-||7.56||7.67||7.65||7.48||7.64 0.13 0.11|
|R16A||7.64||-||7.65||7.64||7.67||7.72||7.72||7.50||7.65 0.07 0.11|
|R39A||7.65||7.76||7.66||7.57||7.52||7.64||7.61||-||7.63 0.08 0.11|
|R49A||7.62||-||-||7.55||7.89||7.56||-||7.42||7.61 0.17 0.11|
|R55A||7.68||7.80||7.66||7.75||7.68||7.73||-||7.62||7.70 0.06 0.11|
|R66||7.61||-||7.56||7.81||7.60||7.60||7.58||7.49||7.61 0.10 0.11|
|R92||7.74||-||7.68||7.78||7.67||7.69||7.66||7.53||7.68 0.08 0.11|
|R102||7.63||-||7.79||7.77||7.61||7.58||7.61||7.47||7.64 0.11 0.14|
|R103||7.69||7.64||7.83||7.59||-||7.64||7.73||7.63||7.68 0.08 0.11|
|R105||-||7.67||7.74||-||7.67||-||7.61||7.61||7.66 0.06 0.11|
|R123||7.55||-||-||7.76||-||7.70||7.68||7.67||7.67 0.08 0.11|
|vB182||7.62||7.65||7.64||7.63||-||-||-||-||7.64 0.02 0.10|
|vB187||7.66||7.67||7.68||7.64||-||-||-||-||7.66 0.02 0.10|
We used the best quality spectra to derive the cluster metallicity. The iron abundance analysis was carried out using MOOG (Sneden 1973 - version December 2000) and Kurucz (1995) model atmospheres. For each star we measured up to eight Fe I lines, whose wavelengths are listed in Table 6 (the EWs are available from S. Randich, upon request); for these lines we adjusted values by carrying out an inverse abundance analysis of the solar spectrum. We used the spectrum of the Sun observed with the FEROS instrument at La Silla, during another observing run. The resolving power of the FEROS spectrum ( ) is somewhat higher than that of our sample spectra; all the lines that we used for the iron analysis, however, do not have close features that could be blended at our resolution and that could lead to overestimate the cluster metallicity. For the Sun, we assumed and the usual solar parameters: K, , and km s-1. Van der Waals broadening was treated using the Unsöld approximation (1955).
The derived iron abundances for each line and the mean abundance for each star are listed in Table 6; the random errors in the mean, and are also listed in the table. The random errors were estimated similarly to Randich et al. (2000); namely, we assumed that, for each star, the standard deviation of the mean iron abundance would be a good estimate of the error - - due to errors in measured EWs and to random uncertainties in atomic parameters, gf-values in particular. We then estimated random errors due to uncertainties in stellar parameters - - by varying each parameter at a time and leaving the other two parameters unchanged; we then quadratically added the related errors. As mentioned above, conservative errors of 100 K in , 0.3 dex in , and 0.3 km s-1 in were assumed. Note that we did not find any abundance trend vs. EW or EP (excitation potential), meaning that our assumed parameters should not be largely in error.
In order to estimate systematic errors and to put our [Fe/H] determination for NGC 6475 on a consistent scale with other well known clusters, we determined [Fe/H] for two Hyades members. The two stars, vB182 and vB187, which have within the range covered by our sample stars, were observed by us during one observing run with UVES on VLT UT2 (the detailed description of the data, analysis, and results will be reported elsewhere). The spectra have resolving power and S/N ratios around 200. We derived stellar parameters for the two Hyades stars consistently with our sample stars and obtained (or assumed in the case of surface gravity): = 5079 K, , km s-1 for vB182 and = 5339 K, , km s-1 for vB187. We measured the EWs of the four Fe I lines of the ones used by us for metallicity included in the UVES spectral range and derived iron abundances as for our sample stars. The abundances for the two Hyades stars are listed in Table 6.
We computed the weighted mean iron abundance for NGC 6475 using all the stars listed in Table 6 and obtained , or . Note that, when computing the mean, we conservatively assumed for each star a total error . The mean for the Hyades is or . In other words a) our metallicity for the Hyades is virtually the same as the usually quoted value for this cluster ( , Boesgaard & Budge 1989), implying that our analysis should not be affected by large systematic errors; b) the metallicity of NGC 6475 is over-solar and very similar to that of the Hyades.
As mentioned in the introduction a metallicity larger than solar
for NGC 6475 was already found by JJ97, who derived
in good agreement with our estimate. In addition, one star
in the sample of JJ97 (JJ26/R14) is
in common with the sample we have used for the
the metallicity we derive for this star is consistent with
the value quoted by JJ97 (
and +0.10, respectively).
However, whereas JJ97 assumed the same
as ours, they assumed a microturbulence
km s-1 for all the stars. Had we assumed this value, we would
have found a much lower metallicity for the cluster (
We note that our choice for the microturbulence parameter
is more in agreement
with other metallicity studies: besides
Boesgaard & Friel (1990), several other authors used low, and
temperature dependent, values of
for cluster and field dwarfs (e.g.,
Edvardsson et al. 1993; King et al. 2000).
In addition, our microturbulence scale is consistent with the
most commonly used value for the solar microturbulence which we also
assumed for our inverse analysis of the solar spectrum (see above).
|Figure 2: Our NGC 6475 equivalent widths are plotted vs. the equivalent widths of JJ97 for the 10 stars in common.|
|Figure 3: NGC 6475 Li abundance vs. : the present sample (filled triangles) is compared to those of JJ97 (circles) and J00 (open triangles); circled symbols denote stars whose membership is still to be confirmed, while down-pointing arrows represent upper limits in (Li). Error bars are also shown.|
|Figure 4: Li abundance vs. for the NGC 6475 merged sample (our sample+JJ97+J00). Top panel: open triangles denote single stars, while filled triangles indicate spectroscopic binary stars. Circled symbols denote stars whose membership is still to be confirmed. Bottom panel: stars with available projected rotational velocity are plotted as open circles: the size of the circles is proportional to ; stars with no CORAVEL information on rotational velocity are plotted as crosses.|
The (Li) vs. distributions for the three samples are plotted in Fig. 3: the three distributions appear very similar; in particular no systematic difference (due, e.g., to different instruments, spectral resolutions or reduction methods) is present; the three sets therefore can be safely merged into a single larger sample. We note that the differences between our EWs and those of JJ97 for the ten stars in common are strongly reduced when considering Li abundances; in the following we will use our own Li measurements for the stars in common.
The Li vs.
distribution of NGC 6475 is almost
flat for the late F stars, showing little depletion
with respect to the initial value ((Li)0=3.1-3.3,
from T Tauri stars and meteorites);
below 6000 K have instead
undergone Li depletion with the Li distribution
showing a rapid decline as the stellar temperature
No evident scatter
in Li abundances is present in this cluster for stars
warmer than 4800 K, although a few stars
below the mean trend are present:
two of these stars have no confirmed membership
= 5239 K and JJ105,
= 5272 K), one is
a probable spectroscopic binary (JJ36,
= 5810 K)
and two stars are confirmed as members
= 5888 K and R51A,
= 5772 K).
The Li spread among cooler stars will be discussed
|Figure 5: Comparison of the Li distributions of the NGC 6475 merged sample (filled triangles) and M 34 (Jones et al. 1997, open circles); circled symbols denote stars whose membership is still to be confirmed.|
In the top panel of Fig. 4 the NGC 6475 merged sample (our stars+JJ97+J00) is shown: single and probable spectroscopic binary stars are plotted with different symbols. This plot shows that, apart from JJ36 (see above) binarity character does not alter significantly the determination of Li abundances because both kinds of stars are well mixed.
In the bottom panel of Fig. 4 NGC 6475 stars with different projected rotational velocities are represented with symbols of different size: there is no evident correlation between Li abundances and values. We stress however that, with exception of two objects with km s-1 (see Sect. 3.1), all the other stars have low rotational velocities.
In Fig. 6 we show a comparison between the Li patterns
NGC 6475, the Pleiades and
We can divide the plot in three temperature ranges:
6000 K: the three clusters
have similar Li distributions, with a mean value slightly
below the meteoritic abundance, i.e. these stars seem to suffer
a very little amount of MS Li depletion.
We do not consider here the Li dip observed in the Hyades stars
6500 K, since our sample contains mostly stars
cooler than the dip.
The only Li poor star belonging to the Hyades
6200 K, but its
membership is uncertain and it is discussed in Thorburn et al. 1993;
K: in this
range, it is evident
that the NGC 6475 Li pattern lies between those of the Pleiades and the Hyades,
suggesting that Li depletion is a continuous process occurring for
both between 100 and 220 Myr and between 220 and 600 Myr.
We stress that below 6000 K a few stars in NGC 6475
appear as depleted as
(or more depleted than) the older Hyades:
these stars seem to be bona fide cluster members
(see Sect. 4.1) and should
be further monitored;
5500 K: as well known, Pleiades stars are characterized
by a large
amount of scatter for late G to late K stars.
The NGC 6475 distribution lies on the lower envelope of the
Pleiades distribution, and several Pleiades stars exist that show the same
amount of depletion as NGC 6475;
is not present in NGC 6475 at least for stars hotter than 4800 K:
more precisely, above this temperature the only two stars
that could indicate the presence of a scatter,
as mentioned, are not confirmed as members;
this shows that at an age of 200-250 Myr, Li abundances have already
converged onto similar values.
This is true also for M 34 (see Fig. 5): note that, whereas our sample
for NGC 6475 is itself statistically significant for stars
hotter than 4800 K,
the two samples together allow us to exclude with an even higher
significance that this result is due to low number statistics.
As already evidenced, M 34 is characterized by a Li scatter among stars cooler
than 4700-4800 K: this dispersion could be present also in NGC 6475,
but the NGC 6475 sample is rather
sparse in this temperature range and we cannot draw any definitive conclusion
about this point.
|Figure 6: The NGC 6475 Li distribution (filled triangles) is compared to those of the Pleiades (S93Jones et al. 1996, circles) and the Hyades (Thorburn et al. 1993, stars). Circled symbols denote stars whose membership is still to be confirmed.|
|Figure 7: Li abundance vs. age for three ranges in ( K, K, K). In the top panel average (Li) values for the Pleiades, NGC 6475+M 34 and the Hyades are shown. Error bars represent standard deviations from the average value; for the Pleiades in the range K, where the Li spread is present, we have plotted the maximum and minimum values of (Li). In the bottom panel theoretical (Li) values at various ages are plotted (D' Antona & Mazzitelli 1994), including two different treatments of overadiabatic convection: MLT is the mixing length theory model, CM is the Canuto & Mazzitelli model.|
A clear effect of [Fe/H] on Li depletion has never been empirically confirmed: for example, the comparison between the Pleiades ( solar) and Blanco 1 ( , Jeffries & James 1999), both with an age of 100 Myr, suggests that PMS Li depletion does not depend on metallicity.
We found NGC 6475 to have ; as mentioned in the introduction, the metallicity of M 34 is somewhat lower, but most likely not as low as found in the early study of Canterna et al. (1979): in fact, Schuler et al. (2003) derived for this cluster . This result is based on five solar-type stars, while, considering their whole sample of nine stars with K, Schuler et al. would have derived ; excluding only the two coolest stars of the total sample, they would have instead found for M 34. The value of the iron content of this cluster is a very crucial point for our discussion: as seen in Fig. 5, there are no significant differences between the (Li) distributions of NGC 6475 and M 34, for stars hotter than 4800 K. The uncertainty on the metallicity of M 34 leads to two different possibilities for the interpretation of this result: (i) if M 34 has (as probable, given the more recent and detailed analysis of Schuler et al. 2003), the similarity between the Li distribution of this cluster and NGC 6475 would not be surprising; (ii) if, on the contrary, M 34 has a lower, close to solar, metallicity, the results of Jeffries & James (1999) based on the comparison of Blanco 1 and the Pleiades would be extended to larger ages, i.e. the overall metallicity does not affect Li depletion up to the age of NGC 6475.
In any case, both possibilities allow us to safely merge the NGC 6475 and M 34 samples to investigate Li evolution as a function of age by comparing these clusters with the younger Pleiades and the older Hyades: in fact, in case (i) we can use the Pleiades in the comparison, since their Li distribution is similar to that of the over-solar metallicity Blanco 1. We will use the Pleiades instead of Blanco 1 since a very rich sample is available for the former cluster, allowing also a more detailed discussion about the spread among K-type stars; if, otherwise, case (ii) is the correct one, the Li patterns of NGC 6475 and M 34 are not affected by metallicity, thus age is the main parameter on which MS Li depletion depends. Finally and obviously, in both cases there is no problem in using the Hyades ( ) in our comparison.
The above discussion is valid for stars warmer than 4800 K; cooler stars deserve a special remark, since these stars in M 34 show a scatter in Li abundances: as mentioned, we cannot draw any definitive conclusion about the presence of a similar spread in NGC 6475. Under case (i), i.e. similar [Fe/H] for the two clusters, one would expect to find a scatter in NGC 6475. If M 34 has instead a solar metallicity (case (ii)) and the scatter exists also in NGC 6475, this would mean that the iron content does not affect Li depletion even for the coolest stars, at least at an age of 220-250 Myr. On the contrary, if further Li observations of NGC 6475 should demonstrate that no spread is present in this cluster, this would suggest that Li depletion in cool stars is affected by metallicity and the mechanism causing the dispersion in Li is also metal dependent.
Whereas we leave the issue of the spread among stars cooler than 4800 K to a future larger sample, we discuss below our results for the hotter stars.
The comparison of the two panels of Fig. 7 shows that: (i) 6000 K- observations agree with the theoretical predictions. The models predict, during both the PMS and MS phases, a temperature at the base of the convective envelope ( ) which is slightly lower than (or at most similar to) the Li burning temperature ( K). The agreement between observations and model predictions allows us to conclude that in late F stars no extra-mixing mechanism is probably present, at least up to the Hyades age; (ii) 5500 K - there is a clear disagreement between theory and observations: the models predict during the PMS, but decreases down to values around before an age of 100 Myr, thus very little Li depletion is present after this age. Standard models cannot explain the observed MS depletion: this confirms that an extra-mixing mechanism is at work in these stars. As suggested by several authors (see Jones et al. 1997 and references therein) extra-mixing could be due to MS angular momentum loss (AML) and angular momentum transport, which, in this case should be a continuous process. If, however, Li depletion is driven by rotational mixing, it is difficult to understand the lack of dispersion among these stars which have, presumably, different rotational histories. (iii) 5000 K - during the PMS, the theoretical is higher than the Li burning temperature, even at very young ages (0.3 Myr), but it decreases after arrival on ZAMS; thus, according to the models, 100 Myr old cluster stars should have depleted a large amount of Li during the PMS.
With regard to the latter point we can consider two opposite hypotheses: a) the upper envelope of the Pleiades is the result of PMS convection only and the lower envelope is over-depleted by the action of an extra-mixing mechanism during PMS; b) the lower envelope of the Pleiades is the result of convection only, while in Li rich Pleiades stars the PMS convection might have been strongly inhibited by some non-standard process; for example, as suggested by Ventura et al. (1998) and D' Antona et al. (2000), the effect of magnetic fields induced by a strong rotation could inhibit Li depletion; alternatively, a strong rotation could significantly modify the stellar structure (see Martín & Claret 1996) and prevent convection and Li depletion: both the proposed processes could explain the dispersion since both rotation and magnetic fields cover a large range of values within the same cluster (e.g. Stauffer et al. 2000).
Case a) appears unlikely since it would imply that, during the PMS, the convection in these cool stars "normally'' do not reach deep enough layers to burn lithium. Hypothesis b) appears more probable; within this hypothesis one can explain both the convergence of Li abundances at the age of NGC 6475 and the large differences in Li depletion time scales in the two intervals [100-220] Myr and [220-600] Myr with the following speculative scenario: in stars of the upper envelope Li depletion is inhibited during PMS (by magnetic fields and/or rotation, see above); then, after the stars have reached the ZAMS (100 Myr), they start loosing angular momentum at a fast rate and extra-mixing occurs, leading to MS Li destruction which is a continuous process from 100 Myr to 600 Myr. Stars on the lower envelope (which are mostly slow rotators) during the PMS deplete a large amount of Li under the action of convection only, which stops at an age of 100 Myr; Li depletion becomes again sensitively efficient around the age of NGC 6475, when the decoupling between the core and the surface is large enough to have extra-mixing due to AML. Note that these cool stars have rather deep convective envelopes and that during MS their remains very close to the temperature which makes Li burning efficient: therefore only a small amount of extra-mixing is required for Li depletion to occur and this explains the large efficiency of Li depletion between 220 and 600 Myr, for stars of both the upper and lower envelopes.
In summary, we suggest that the spread observed among Pleiades stars cooler than 5500 K could be due to the fact that convection and Li depletion may be inhibited by processes related to rotation and/or magnetic fields, which vary from star to star; we also conclude that the convergence of Li abundances at the age of NGC 6475 for stars hotter than 4700-4800 K could be due instead to extra-mixing mechanisms, which drive the depletion after arrival on the ZAMS, and have different time scales depending on the initial rotation. Finally, we stress that in the discussion above we assumed that the spread observed in the Pleiades and M 34 is due to a real scatter in Li abundances. We mention that several authors suggested that the spread in Li equivalent widths could not necessarily correspond to a real spread in abundances, and they investigated whether the scatter could be due to the effects of surface activity (spots in particular) on the line formation and strength (e.g. Stuik et al. 1997; King et al. 2000; Randich 2001; Barrado y Navascués et al. 2001). The issue of the scatter in Li abundances remains therefore open.
Our main results are:
a) We confirm the over-solar metallicity of the cluster; specifically we found .
b) The comparison of NGC 6475 with the similar age M 34 shows no significant differences between the two Li distributions down to 4700-4800 K. This is not surprising, given the small difference in metallicity between the two clusters, according to Schuler et al. (2003) which found for M 34. M 34 shows an evident Li abundance spread among stars cooler than 4700-4800 K; the NGC 6475 sample is instead rather sparse in this temperature range. Thus, although there may be an indication for the presence of a dispersion, no definitive conclusion can be drawn. More late K-type stars in NGC 6475 should be observed.
c) Assuming that metallicity does not affect PMS Li depletion, as shown by the Pleiades-Blanco 1 comparison, we can consider the age sequence from the Pleiades, to NGC 6475+M 34 and then to the Hyades. We found that Li depletion occurs during the MS phase of G and K-type stars; the Li pattern of NGC 6475 for both G and early K stars lies between those of the Pleiades and of the Hyades. This means that extra-mixing mechanisms are likely at work both between 100 and 220 Myr and between 220 and 600 Myr.
d) The star-to-star scatter in Li abundance observed among stars cooler than 5500 K in clusters as young as (and younger than) the Pleiades is not present in NGC 6475 stars hotter than 4700-4800 K, as well as in M 34 stars over the same temperature range. We suggest that the spread observed in the Pleiades could be due to processes related to rotation and magnetic fields, which inhibit convective mixing and Li depletion during the PMS for part of the stars (those in the upper envelope); the disappearing of the scatter at the age of NGC 6475 (for stars in the temperature range [5500-4700] K) could be due to extra-mixing processes, which could also be responsible for the acceleration of Li depletion between this age and that of the Hyades.
As a final remark, we stress that the determination of oxygen and other elements abundances are a very important issue for the investigation of Li evolution.
This work has partially been supported through a grant by Ministero dell'Istruzione, Università e Ricerca (MIUR) to S. Randich and R. Pallavicini. We thank the referee, Dr. B. F. Jones, for his very useful comments.