Free Access
Issue
A&A
Volume 535, November 2011
Article Number A30
Number of page(s) 21
Section Galactic structure, stellar clusters and populations
DOI https://doi.org/10.1051/0004-6361/201117473
Published online 27 October 2011

© ESO, 2011

1. Introduction

Open clusters (OC) are ideal test particles for studying the evolution of metallicity with time, inferring the so-called age-metallicity relation, and with Galactocentric radius, the metallicity gradient, measuring in the Galactic disc. Their properties can be determined with smaller uncertainties than for field stars, since they are coeval group of stars at the same distance that have a homogeneous chemical composition. Unfortunately, of the  ≃1700 known OC (e.g., Dias et al. 2002), only  ≃140 possess some metallicity determination, mostly obtained from photometric indicators, such as Washington or Strömgren photometry (see Twarog et al. 1997; Chen et al. 2003, and references therein) and low-resolution spectroscopy (e.g., Friel & Janes 1993; Friel et al. 2002).

The most accurate way to determine the chemical abundances is to analyse high-resolution spectroscopy. It allows us to investigate not only metallicity, but also abundance ratios – with respect to iron or hydrogen – of other chemical species such as α-elements, s-process elements, and r-process elements, which are synthesised in different environments and on different timescales (e.g., SNe Ia, SNe II, giants, supergiants, etc). In the past few years, a number of research groups have addressed the challenge of increasing the number of OC with chemical abundances determined from high-resolution spectroscopy (e.g., Sestito et al. 2004; D’Orazi et al. 2006; Sestito et al. 2006; Bragaglia et al. 2008; Pace et al. 2008; D’Orazi et al. 2009; Friel et al. 2010; Pace et al. 2010; Pancino et al. 2010a; Jacobson et al. 2011a). However, the number of OC with chemical abundances determined with this technique is still small (see Sect. 5), and significant uncertainties remain in the determinations of both the metallicity gradient and the age-metallicity relation, which are the fundamental ingredients of chemical evolution models.

Table 1

Observing logs and programme star properties.

In this paper, the second of a series initiated by Pancino et al. (2010a, hereafter Paper I), we present high quality and homogeneous measurements of chemical abundances for red clump stars in four OC: Be 32, NGC 752, Hyades, and Praesepe. The Hyades is the nearest OC and its four known red giants have been widely studied (Schuler et al. 2009; Mishenina et al. 2007; Fulbright et al. 2007; Schuler et al. 2006; Mishenina et al. 2006; Boyarchuk et al. 2000; Luck & Challener 1995), hence it provides a very good reference frame to compare our abundances with the literature. Both NGC 752 and Praesepe have been well-studied, but all information about their chemical composition is based mainly on their main-sequence stars (e.g., Pace et al. 2008; An et al. 2007; Sestito et al. 2004; Burkhart & Coupry 1998; Hobbs & Thorburn 1992). To our knowledge, there have been no recent measurements of the chemical abundance of their giants from high-resolution spectroscopy. Finally, Be 32 has been the subject of some studies (e.g., Richtler & Sagar 2001; Friel et al. 2010; Bragaglia et al. 2008; D’Orazi et al. 2006). The properties and previous studies of each cluster is described in more depth in Sect. 4.

This paper is structured as follows: observations and data reduction are described in Sect. 2; equivalent-width measurements are presented in Sect. 3, together with the abundance analysis and its uncertainties; results are compared with the literature in Sects. 46; and finally our main conclusions are summarised in Sect. 7.

2. Observational material

A total of twelve stars spread in the four OC were observed. They were selected from the WEBDA1 database (Mermilliod 1995), and the 2MASS2 survey (Cutri et al. 2003; Skrutskie et al. 2006). Table 1 summarizes the identifications, coordinates, and magnitudes of each target star. Their position in the color − magnitude diagram taken from D’Orazi et al. (2006), Johnson (1953), Johnson & Knuckles (1955), and Johnson (1952) for Berkeley 32, NGC 752, Hyades, and Preasepe, respectively, are shown in Fig. 1.

Observations were carried out with the fibre echelle spectrograph FOCES (Pfeiffer et al. 1998) attached at the 2.2 m Calar Alto Telescope (Almeria, Spain) between the 1 and 3 of January 2005. The chosen set-up provides a spectral resolution (R = λ/δλ) of about 30 000. In summary, all stars were observed in 2 − 7 exposures lasting 10 − 30 min each, depending on their magnitudes, until a global signal-to-noise ratio (S/N) of at least 60 per pixel was reached around 6000 Å. Exposures with S/N  <  20 were neglected because they were too noisy. Finally, sky exposures as long as our longest exposures (30 min) were taken, but the levels were sufficiently low for us to avoid sky subtraction (as in Paper I). The number of useful exposures, the total integration time, and the global S/N for each star are listed in the last three columns of Table 1.

thumbnail Fig. 1

Location of target stars (large black dots with star ID labels) in the color–magnitude diagrams of their respective parent clusters (small grey dots).

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2.1. Data reduction

Various steps of data reduction were performed exactly as in Paper I. Briefly, the frames were de-trended with the IRAF3 tasks ccdproc and apflatten. The spectra were then extracted, wavelength-calibrated, normalized, and the echelle orders were merged using tasks in the IRAF echelle package. Finally, the noisy ends of each combined spectrum were cut, allowing for an effective wavelength coverage from 5000 to 9000 Å.

Before combining all exposures of each star, we removed sky absorption features (telluric bands of O2 and H2O) with the help of the IRAF task telluric. The same two hot, rapidly rotating stars, HR 3982 and HR 8762, of Paper I were used. The strong O2 band around 7600 Å had been saturated and therefore could not be properly removed. This spectral region was not used for the abundance analysis, in addition to the small gaps between echelle orders that appeared after λ  ≃  8400 Å.

2.2. Radial velocities

Table 2

Heliocentric radial velocity measurements and 1σ errors (Vr ± δVr)here for each programme star.

We used DAOSPEC (Stetson & Pancino 2008) to measure the observed radial velocities for each individual exposure with S/N  ≥  20, using  ≃ 300 absorption lines of different elements, with typical uncertainties of about 0.1 km s-1 (see Paper I for details). We used the same linelist as the one used for abundance determinations (see Sect. 3 for details). Heliocentric corrections were obtained with the IRAF task rvcorrect, with a negligible uncertainty of smaller than 0.005 km s-1. We also used DAOSPEC to determine the absolute zero-point of the radial velocity determinations, using a list of telluric absorption lines as the input linelist, obtained from the GEISA4 database (Jacquinet-Husson et al. 1999, 2005). The resulting zero-point corrections, based on  ≃250 telluric lines, are generally no larger than  ± 1 km s-1, with a typical error of about  ≃ 0.5 km s-1.

The final values, computed as the weighted mean of heliocentric velocities resulting from each exposure of the same star, are listed in Table 2. Our determinations are generally in close agreement with literature values to within 3σ, except for star 208 in NGC 752, which has a slightly smaller radial velocity than other objects in this cluster. The fact that this star was recognised as a spectroscopic binary (see Pourbaix et al. 2004; Mermilliod et al. 2007) explains the disagreement. According to its radial velocity curve (Mermilliod et al. 2007), we observed this binary near minimum, which implies that we observed only one of the components of the system. For this reason, and because derived abundances are in good agreement with those of other stars in the same cluster, we retained this object in our final sample. In summary, we considered all the observed targets as likely members of their respective clusters.

2.3. Photometric parameters

First guesses of the atmospheric parameters effective temperature (Teff), logarithmic gravity (log g), and microturbulent velocity (vt), for our target stars were derived from a photometric data listed in Table 1, as described in Paper I. In brief, Teff were obtained using the Alonso et al. (1999) and Montegriffo et al. (1998) colour − temperature relations, both theoretical and empirical, and the dereddened colours (B − V)0, (V − IJ)0, (V − R)0, and (V − KS)0. We assumed the E(B − V) values listed in Table 3 and the reddening laws of Cardelli et al. (1989). In the case of Be 32, we have IC magnitudes instead of IJ ones, so we dereddened (V − IC) with the law of Dean et al. (1978), and converted it into (V − IJ)0 with the transformations by Bessell (1979). The 1σ errors in each Teff estimate were computed using the magnitude and reddening uncertainties together with the standard deviation in the colour − temperature relationships used. The photometric Teff estimates, listed in Table 4, are the weighted mean of the different values obtained from each considered colour and colour − temperature relations.

Photometric gravity estimates were derived from the above Teff and the bolometric corrections, BCV, derived using the Alonso et al. (1999) prescriptions and the fundamental relationships \arraycolsep1.75ptwhere red clump masses, listed in the last column of Table 4, were extrapolated from Table 1 of Girardi & Salaris (2001). We assumed that log g = 4.437, Teff,⊙ = 5770   K and Mbol,⊙ = 4.75, in conformity with the IAU recommendations (Andersen 1999). As above, we averaged all our estimates to obtain log g(phot), listed in Col. 5 of Table 4.

As discussed in Paper I, the photometric estimate of the microturbulent velocity, vt, was obtained using the prescriptions both of Ramírez & Cohen (2003), vt = 4.08−5.01  ×  10-4   Teff, and of Carretta et al. (2004), vt = 1.5−0.13   log g. The latter velocity, which takes into account the effect described by Magain (1984)5, is on average lower by Δvt = 0.50  ±  0.03 km s -1 than the Ramírez & Cohen (2003) estimate. Therefore, we chose not to average the two estimates, but to use them as an indication of the vt range to explore in our abundance analysis (see Sect. 3.2).

Table 3

Adopted cluster parameters.

Table 4

Stellar atmosphere parameters for the programme stars (see text).

Table 5

Equivalent widths and atomic data of the programme stars.

3. Equivalent widths and abundance analysis

We used the same linelist as that described in Paper I. In brief, all lines and their atomic data were extracted from the VALD6 database (Kupka et al. 1999), with a few exceptions (see Paper I for details). Briefly, for some highly discrepant Mg lines, we used the NIST loggf values; we used the Johansson et al. (2003) loggf for the Ni line that contaminates the [O I] line at 6300 Å, and provides oxygen abundances more in line with the other α-elements; we used the Nd loggf values by Den Hartog et al. (2003), which minimize the spread in the Nd abundance. Finally, we tried both the VALD and the NIST values for Ca, finding an average difference of 0.17 dex (see Paper I). There is no special reason for choosing NIST over VALD (or vice-versa), so we kept the VALD values to help maintain some homogeneity, but we note that the Ca loggf values carry a large uncertainty of the order of 0.2 dex.

3.1. Equivalent widths with DAOSPEC

The task DAOSPEC (Stetson & Pancino 2008) was used to automatically find and measure equivalent widths (EW), by performing a Gaussian fitting of the identified lines. DAOSPEC provides a formal error in the Gaussian fit, δEW, and a quality parameter, Q (see Stetson & Pancino 2008; and Paper I, for more details). The relative error δEW/EW and the quality parameter Q can be used to distinguish good and bad lines, and they were indeed used to select the highest quality lines for the abundance analysis, as described in detail in Paper I. The measured EW for our program stars are shown in the electronic version of Table 5 along with the δEW and Q parameter estimated by DAOSPEC.

Four of our target stars have published EW measurements from high-resolution spectra. These consist of three stars (namely, 028, 041, and 070) observed in the Hyades by Boyarchuk et al. (2000) with R ~ 45   000, and star 0456 in Be 32 studied by Bragaglia et al. (2008) with R ~ 40   000. We have a total of 100, 92, and 51 lines in common for stars 028, 041, and 070 in the Hyades, respectively, and 51 lines for star 0456 in Be 32. Figure 2 compares the comparison between the EW determined with DAOSPEC with the values published by Bragaglia et al. (2008) and Boyarchuk et al. (2000). The differences (see Fig. 2) are negligible within the uncertainties; we find a small offset of 5.6 mÅ in the case of star 041 in the Hyades, which is however still within 1σ. We can therefore consider our measurements in good agreement with similar studies.

thumbnail Fig. 2

Comparison of our EW measurements with those by Bragaglia et al. (2008) for star 0456 in Be 32, and by Boyarchuk et al. (2000) for three Hyades giants. Dotted lines mark perfect agreement (zero difference), while dashed lines are linear fits to the data.

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3.2. Abundance analysis

Abundance calculations and spectral synthesis (for oxygen) were performed using the latest version of the abundance calculation code originally described by Spite (1967). We used the MARCS model atmospheres developed by Edvardsson et al. (1993). We also used of ABOMAN, a tool developed by E. Rossetti at the INAF, Bologna Observatory, Italy, which allows the semi-automatic processing of data for several objects, using the aforementioned abundance calculation code. The tool ABOMAN performs all the steps needed to choose the best-fit model automatically (see below) and compute abundance ratios for all elements, and provides all the graphical tools required to analyse the results.

The detailed procedure followed to derive the chemical abundances is described in depth in Paper I. In brief, we calculated Fe I and Fe II abundances for a set of models with parameters extending  ± 3σ around the photometric estimates of Table 4. We chose the model that satisfied simultaneously the following conditions: (i) the abundance of Fe I lines should not vary with excitation potential χex; (ii) the abundance of Fe I lines should not vary significantly with EW, i.e., strong and weak lines should infer the same abundance7; (iii) the abundance of Fe I lines should not differ significantly from the abundance of Fe II lines; and (iv) the abundance of Fe I lines should not vary significantly with wavelength.

Once the best-fit model has been found, abundance ratios of all the measured elements were determined, as shown in Table 6, as the average of abundances given by single lines. The internal (random) errors were then computed as σ/. Oxygen abundances were determined by means of spectral synthesis of the region around the [O I] forbidden line at 6300 Å. In this case, the internal uncertainty was estimated using the average abundance difference between the best-fit spectrum and two spectra placed approximately 1σ (of the Poissonian noise) above and below it. Average cluster abundances (Tables 7) were computed as weighted averages of abundance ratios of single stars.

Comparison of our results with available literature is discussed in details in Sect. 4.

3.3. Abundance uncertainties and the Sun

The internal (random) uncertainty described above includes uncertainties related to the measurement of EW and to the atomic parameters (dominated by loggf determinations). We must consider other sources of uncertainty (see Paper I for details) such as: the uncertainty owing to the choice of atmospheric parameters; the uncertainty owing to the continuum normalization procedure; the uncertainty in the reference solar abundance values.

Uncertainties due to the choice of stellar parameters were evaluated with the method proposed by Cayrel et al. (2004). In brief, we altered the predominant atmospheric parameter, i.e., by altering one atmospheric parameter, Teff, within its uncertainty (~100 K) and re-optimizing the other parameters for the hottest and coolest stars in our sample. We re-calculated abundances with the procedure described in the previous Section. The external uncertainties, listed in the last column of Table 6, are estimated by averaging errors calculated with the higher and lower temperatures for the warmest and coolest stars in our sample (namely, stars 001 and 208 in NGC 752).

Uncertainties due to the continuum normalization procedure might also affect the obtained EW and, therefore, the derived abundances. Their contribution is estimated by averaging the differences between the EW obtained with the “best-fit” continuum and those derived by lowering and raising the continuum level by the continuum placement uncertainty. This is calculated from Eq. (7) of Stetson & Pancino (2008). The typical uncertainty caused by the continuum placement is ΔEW  ~  1 mÅ and almost independent of the EW. This small uncertainty has a negligible impact on the derived abundances in comparison with other sources of uncertainty described above. Therefore, they have not been explicitly included in the error budget.

To validate the whole procedure used here, in Paper I we performed an abundance analysis of the ESO HARPS solar spectrum reflected by Ganymede8. We used the same line list, model atmospheres, and abundance calculation code that we used on our OC target stars, and found solar values for all elements, with the only marginal exceptions of barium and aluminium (see also Sect. 5). While the details of this analysis can be found in Paper I, we mention here that our reference solar abundances are taken from Grevesse et al. (1996).

4. Cluster-by-cluster discussion

4.1. Berkeley 32

Berkeley 32 (α2000 = 06h58m07s and δ2000 =  + 06°25′43′′) is a distant OC (Rgc = 11.6 kpc) located towards the Galactic anticentre and situated 260 pc above the disc plane. Its distance makes it one of the crucial clusters for a correct determination of the metallicity gradient along the Galactic disc, and therefore one of the key OC to the understanding of disc formation and evolution. The color − magnitude diagram of this cluster (e.g., D’Orazi et al. 2006), contaminated by disc stars, shows a clear main sequence turn-off with a sparsely populated red giant branch. Determinations of its age, mainly using morphological indicators, yield a value of  ≃ 5 Gyr (e.g., D’Orazi et al. 2006; Salaris et al. 2004; Richtler & Sagar 2001; Carraro & Chiosi 1994; Kaluzny & Mazur 1991).

Given its large distance, it has not been well-studied spectroscopically, but we could compare our results with two recent high-resolution studies of Bragaglia et al. (2008) and Friel et al. (2010). We found a very close agreement of our abundance ratios with those studies (see Table 8). The exceptions are Ba and Na. It is well-known that Ba abundances are enhanced by HFS (e.g., D’Orazi et al. 2009) effects that should explain the differences from Bragaglia et al. (2008). The [Na/Fe] ratio is lower than the values reported by Bragaglia et al. (2008) and Friel et al. (2010) by  − 0.25 and  − 0.32 dex, respectively. The difficulty in measuring Na lines, which suffer from NLTE effects, could easily explain this controversy. Moreover, different model atmospheres, stellar and atomic parameters, etc., between different studies may also play a role (and remove this discrepancy).

Table 6

Abundance ratios for single cluster stars, with their internal and external (last column) uncertainties.

Table 7

Average cluster abundances, obtained as the weighted average of the single stars abundances in each of them.

4.2. NGC 752

NGC 752 (α2000 = 01h57m41s, δ2000 =  + 37°47′06′′) is an old (~1.6 Gyr) OC located in the solar neighbourhood at a distance of  ≃ 400 pc. This cluster has a low central concentration and contains a relatively small number of members. Its color − magnitude diagram (e.g., Johnson 1953) has a still poorly understood morphology. The turn-off area is well-populated by early F-type stars, while the low main-sequence appears to be sparsely populated (Fig. 1). This, together with the age of this cluster, may be an indication of the dynamic escape of low mass stars. Stellar evolution models also predict a well-populated red giant branch, which is not observed. All the known red giants are located in the red clump region (Bartašiūtė et al. 2007), which has a peculiar morphology because it has a faint extension slightly to the blue of its main concentration, which cannot be reproduced by stellar evolution models (Girardi et al. 2000).

Photometry and low/medium resolution spectroscopy studies (see Bartašiūtė et al. 2007, and references therein) determined a slightly subsolar metallicity (i.e. [Fe/H] =  − 0.16  ±  0.05, Friel & Janes 1993). A similar result was found with high-resolution spectroscopy (R  ≃  40 000, S/N  ≃  80 − 150) in eight F-type stars around the main sequence turn-off ([Fe/H] =  − 0.09  ±  0.05, Hobbs & Thorburn 1992). However, an investigation based on high-resolution spectroscopy (R  ≃  57 000, S/N  ≃  30 − 80) of 18 G giant stars obtained a solar [Fe/H] ratio ([Fe/H] = +0.01  ±  0.04, Sestito et al. 2004) in closer agreement with the value determined here. To our knowledge, we are the first to publish abundance ratios of elements other than [Fe/H] for this cluster.

4.3. Hyades

The Hyades cluster (Melotte 25; α2000 = 04h26m54s and δ2000 =  + 15°52′00′′) is the closest OC to the Sun (~45 pc) located in the constellation of Taurus. Its proximity has motivated an extensive study lasting more than a century (starting with Hertzsprung 1909). The OC is embedded into a moving group with the same name, which suggests that it would be part of a dynamical stream coming from the inner Galaxy or a disrupting cluster (Famaey et al. 2007).

Table 8

High-resolution average Be 32 abundances.

Being one of the most studied clusters, both photometrically and spectroscopically, it is the ideal cluster for abundance analysis comparisons. The color − magnitude diagram of this young OC (~0.7 Gyr, see Table 3; e.g., Johnson & Knuckles 1955) contains only four red giant stars that have been confirmed as members from their parallaxes, proper motions, and radial velocities. Most of the existing abundance studies are focused on main sequence stars (see e.g., Paulson et al. 2003; Burkhart & Coupry 2000; Varenne & Monier 1999,and references therein). A comparison of the Hyades average abundances determined from some (or all) of the known four red giants are shown in Table 10. The averages of the abundances compiled until 1999 by Varenne & Monier (1999) are shown in the last column of Table 10 for reference. In general, [Fe/H] appears slightly supersolar, while all other abundance ratios are solar, and our abundance ratios agree well with literature values.

The three late-type Hyades giants (028, 041, and 070) have been widely studied (e.g., Luck & Challener 1995; Boyarchuk et al. 2000; Schuler et al. 2006, 2009; Mishenina et al. 2006, 2007; Fulbright et al. 2007). In Table 9 we compiled available literature data. Our temperatures are slightly lower (by  ~100 K) than the literature ones, whereas our values of log g and vt are similar. These marginal differences appear to have no significant impact on the derived abundance ratios, which agree very well with literature ones. Exceptions are Al, Ba, and O, which suffer from technical measurement problems (not strictly related to the Hyades cluster) and are discussed in Sects. 3.3 and 5.

4.4. Praesepe (NGC 2632)

The cluster popularly known as Praesepe or Beehive (also called M 44, NGC 2632 or Melotte 88; α2000 = 08h40m24s and δ2000 =  + 19°40′00′′ ), is an old OC (0.65 Gyr, see Table 3) well known from the antiquity. It is located in the Cancer constellation at a distance of  ≃ 175 pc, computed from Hipparcos parallaxes.

Its metal content was derived with different methods (e.g., Friel & Boesgaard 1992; Komarov & Basak 1993; Claria et al. 1996; Hui-Bon-Hoa & Alecian 1998; Burkhart & Coupry 1998, 2000; Dias et al. 2002; Pace et al. 2008). In general, all the quoted studies obtained a metallicity either barely or definitely supersolar. Of these, the high-resolution abundance determinations were derived mainly for dwarfs or early-type giants (e.g., Friel & Boesgaard 1992; Burkhart & Coupry 1998; An et al. 2007; Pace et al. 2008). Surprisingly, to our knowledge, there are no recent high-resolution abundance determinations of late-type giants in this cluster.

Table 9

Abundance comparison of individual Hyades stars (see text).

Table 11 shows a comparison of our results with some of the most recent high-resolution studies. In general, the [Fe/H] we derived in our late-type giants lies in-between those of Pace et al. (2008) and An et al. (2007), suggesting that the proposed dichotomy of literature values (barely supersolar versus definitely supersolar) should be interpreted rather as an above average uncertainty. This larger than usual uncertainty could naturally arise from the different spectral types and abundance analysis methods employed in the literature. The [Fe/H] ratio derived by (Burkhart & Coupry 1998), based on Am stars, is on average  ≃ 0.3 dex larger than the values obtained in other works using different spectral-type stars. Although these stars should in principle reflect the chemical composition of the cluster, Am stars always have overabundant Fe abundances relative to other objects in the same clusters, without a clear explanation appearing in the literature. As in the case of the Hyades, Na and O abundances derived by us appear marginally discrepant with those by Pace et al. (2008), and will be discussed in more detail in Sect. 5.

5. Discussion of abundance ratios

As in Paper I, we compared our abundance ratios (and those from Paper I) with both others in the literature and the abundances of the Galactic disc field stars from Reddy et al. (2006, 2003) in Figs. 3 to 6. We extended the open cluster abundance compilation of Paper I (see Table 12) with both recent published works and old studies that were not included in the previous version. In both cases, as in Paper I, we included only studies based on high-resolution (R  ≳  18 000) spectroscopy. When more than one determination was available for one cluster, we simply plotted them all to give a realistic idea of the uncertainties involved in the compilation, and we did not attempt to correct for differences between the abundance analysis procedures (loggf, solar reference, and so on), because this would be beyond the scope of the present article.

5.1. Iron-peak element ratios

Figure 3 shows the abundance ratios of iron-peak elements. Our OC with abundances close to solar (i.e., Hyades, Praesepe, and NGC 752) are in very good agreement with the results obtained in other OC studied with high-resolution spectra and in disc stars of similar metallicity. A larger scatter or marginal discrepancies are sometimes observed for the odd elements Sc, V, and Co, but this is because of the well-known hyperfine structure (HFS) of the lines usually employed in the analysis. The element that appears to suffer more from these effects is vanadium. This scatter is due, at least in part, to the different procedures used in the literature for treating the HFS splitting. We stress that in our case, we do not attempt any HFS correction.

The most metal-poor and oldest OC in our sample, Be 32, has a puzzling behaviour. While all its iron-peak abundance ratios are still compatible with the literature values for OC and field stars of similar metallicity (uncertainties are large), some underlying discrepancy could be present. For example, HFS should cause an overestimate (and not an underestimate) of vanadium. In addition, chromium appears to be lower than solar. We note that (see Table 8) the literature Co and Cr determinations by Friel et al. (2010), Sestito et al. (2006), and Bragaglia et al. (2008) are very similar to ours. In the case of our [Cr/Fe] measurement for Be 32, we must note that our two giants appear to exhibit quite different [Cr/Fe] abundances, resulting in a large scatter in the cluster average value. This large scatter is most probably due to a measurement uncertainty, and should not be considered significant.

Table 10

High-resolution average Hyades abundances from giants.

Table 11

High-resolution average Praesepe (NGC 2632) abundances.

thumbnail Fig. 3

Comparison between our iron-peak abundance ratios (large black dots), those of Paper I (large black open circles), high-resolution measurements listed in Table 12 (large dark grey dots), field stars belonging to the thin disc (light grey dots, Reddy et al. 2003), and to the thick disc (tiny light grey dots, Reddy et al. 2006). Errorbars in our measurements are the quadratic sum of all uncertainties discussed in Sects. 3.2 and 3.3.

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In Paper I, we noticed a peculiar behaviour in the Ni abundance ratios of literature OC determinations: they appear to be slightly richer in Ni than field stars by roughly 0.05 dex. Our [Ni/Fe] ratios are in closer agreement with the field star determinations than with the OC ones. Although this difference is small (within the uncertainties), it appears systematic in nature, and we were unable to find any easy explanation, such as the choice of either solar reference abundances or the loggf system, of this discrepancy.

5.2. Alpha-element ratios

thumbnail Fig. 4

Comparison of our α-elements ratios with literature values. Symbols are the same as in Fig. 3.

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thumbnail Fig. 5

Comparison of our s-process elements ratios with the literature ones. Symbols are the same as in Fig. 3, except for the black star-like symbols in the top [Ba/Fe] panel, which represent the revision of Ba abundances with spectral synthesis performed by D’Orazi et al. (2009).

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Figure 4 shows the abundance ratios of α-elements. As for iron-peak elements, our measurements are always compatible with the literature values, within their uncertainties. Generally speaking, all our OC show roughly solar α-enhancements, even Be 32, which has a lower metallicity.

However, some elements deserve some more discussion, as was noted in Paper I. For example, the loggf of calcium are quite uncertain, and we chose the VALD reference atomic data, which explains why our [Ca/Fe] ratios are slightly lower than the bulk of literature determinations for cluster and disc stars. A similar problem affects the Mg lines, as can clearly be appreciated from the large spread of literature values. Our [Mg/Fe] determinations tend to lie on the upper envelope of literature ratios for OC. A deeper discussion of Mg abundances can be found in Paper I.

In the case of oxygen, the problem is instead in the difficulty in measuring its small lines. The forbidden [O I] line at 6300 Å, which we used in this paper, suffers from contamination by a Ni line and by telluric absorption features, while the O triplet around 7770 Å (used by some other studies) suffers from NLTE effects. This is reflected by the large scatter in literature values.

5.3. Heavy element ratios

We determined abundances for three heavy s-process elements: Ba, La, and Nd; and one light s-process element: Y (Fig. 5). Literature determinations for these elements are not numerous. D’Orazi et al. (2009) measured Ba in several OC using spectral synthesis to take into account HFS. The [Ba/Fe] abundances derived by D’Orazi et al. (2009) taking into account HFS do not differ significantly from other literature determinations (including ours). The [Ba/Fe] ratios are clearly above solar for most of the clusters and they show a scatter larger than  ~0.5. D’Orazi et al. (2009) found this scatter to be due to age: the Ba content appears to increase for younger clusters. The available lanthanum and neodymium lines were unfortunately relatively small, and we were able to find fewer published studies to compare with. As a result, the solar clusters (Hyades, Praesepe, and NGC 752) have La and Nd ratios in good agreement with the literature, while Be 32 appears to have lower [La/Fe] and [Nd/Fe] than the few studied OC at a similar metallicity, which are Mel 66 (Gratton & Contarini 1994) and NGC 2243 (Smith & Suntzeff 1987). However, our [Nd/Fe] agrees well with the field star solar ratios. The only light s-process element we could measure, Y, relies on a couple of weak lines that provide uncertain abundances (see the large errorbar in Fig. 5). Our Y ratio appears to be lower than all literature estimates, although still compatible with the solar values of field stars of similar metallicity, within the large uncertainties.

In summary, we can say that all the studied clusters appear to have roughly solar s-process enhancements, but it would be extremely interesting to attempt a more detailed study of s-process elements in OC, as done by D’Orazi et al. (2009) for barium.

5.4. Ratios of Na and Al and anticorrelations

thumbnail Fig. 6

Comparison between our [Na/Fe] and [Al/Fe] ratios and the literature ones. Symbols are the same as in Fig. 3.

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As discussed in Paper I, the study of light elements in OC is quite interesting. The elements Al and Na, together with Mg, O, C, and N, show puzzling (anti-)correlations in almost all of the studied globular clusters, in the Milky Way (see, e.g., Carretta et al. 2010; Pancino et al. 2010b, and references therein) and outside (e.g., Mucciarelli et al. 2009; Letarte et al. 2006). No (anti-)correlations were observed in either field stars (but see Martell & Grebel 2010) or OC (Martell & Smith 2009; de Silva et al. 2009; Smiljanic et al. 2009; and Paper I) so far. This suggests that metallicity, cluster size and age, or the environment must play a rôle, and therefore finding (anti-)correlations in some OC would be of enormous importance to put tighter constraints on the phenomenon.

We determined abundances of Al and Na and compared them with published results in Fig. 6. While in the case of aluminium the agreement with literature values is good, we find a significantly lower [Na/Fe] ratio for Be 32 than for other clusters or field stars of similar metallicity. Generally speaking, the large scatter in Na determinations could be due to the difficulties in measuring Na lines, often affected by NLTE effects (Gratton et al. 1999), although no such scatter is observed among field stars. However, a few clusters have [Na/Fe] lower than our Be 32 determination, and NLTE corrections (Gratton et al. 1999) could make the discrepancy of our Be 32 Na determination even worse. Unfortunately, given the large scatter and the difficulty of measurement, it is difficult to either confirm or exclude the presence of some (small) intrinsic [Na/Fe] scatter in this clusters.

In Fig. 7, all the studied stars occupy the “normal stars” loci, which is around solar for Na and Al, and slightly α-enhanced for O and Mg (see Sect. 5.2). There is a hint of correlation between [Al/Fe] and [Na/Fe], which was also observed for objects studied in Paper I. Of course, small variations in Teff could induce artificial correlations between element pairs, so the observed trend is most probably not-significant. However, we again note that the Na spread is very large (see above), suggesting that a small degree of chemical anomalies (barely hidden within the present observational uncertainties) cannot be completely excluded.

thumbnail Fig. 7

A search for (anti)-correlations of Al, Mg, Na, and O among our target stars. The four panels show different planes of abundance ratios, where stars belonging to each cluster are marked with different symbols. Dotted lines show solar values, solid lines show linear regressions and the typical uncertainty (~0.1 dex) is marked at the lower right corner of each panel.

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6. Galactic trends

The existence of trends in the chemical abundances with Galactocentric distance, Rgc, vertical distance to the Galactic plane, z, and age, are key to understanding Galactic disc formation and evolution because they provide fundamental constraints on chemical evolution models. Different tracers have been used to investigate trends in the Galactic disc: OB stars (e.g., Daflon & Cunha 2004), Cepheids (e.g., Lemasle et al. 2008), H II regions (e.g., Deharveng et al. 2000), and planetary nebulae (e.g., Costa et al. 2004). However, as coeval groups of stars at the same distance and with a homogeneous chemical composition, OC are the ideal test particles to investigate the existence of radial and vertical gradients and of an age-metallicity relation in the disc.

We complement the small sample of abundance ratios obtained here and in Paper I with a revised version of the literature data first presented in Paper I (Table 12). When a cluster had two or more abundance determinations available in the literature, we averaged them to make the figures easily readable and the error bars are, simply, calculated as the standard deviation. For those clusters with only one abundance determination, the error bars are the uncertainties in those determinations. The heliocentric distances compiled in the updated version of the Dias et al. (2002) database were used to obtain Rgc and z for each cluster, assuming RGC = 8.5 kpc. Ages were obtained from the same source, which is a compilation of different values available in the literature, hence might still be quite inaccurate for some clusters. In spite of its heterogeneity, our compilation contains a total of 89 clusters and is, to the best of our knowledge, the largest available in the literature, based on high-resolution spectroscopic abundances. Any attempt to homogenize this sample, for which abundances, distances, and ages have been derived from very different techniques, is clearly beyond the scope of this paper. This prevents us from a detailed analysis of the Galactic trends of all elements. For this reason, we focus only on [Fe/H] and [α/Fe] ratios. In spite of this heterogeneity, this analysis is still very useful owing to the number of clusters, and the large range of ages, and vertical and radial distances covered, even if the heterogeneity of the sample forces us to be extremely cautions when drawing any conclusion from the data.

6.1. Trends with Galactocentric radius

Radial gradients may arise when the disc forms, and different mechanisms can produce them: for example, different timescales of star formation at different distances (e.g., Schaye 2004); a radial variation in the infall of gas; or a change in the yield as a function of the radius (e.g., Molla et al. 1996). This initial radial gradient can be either amplified (steepened) or washed out (flattened) with time by radial mixing (e.g., Roškar et al. 2008).

Since the pioneering work of Janes (1979), OC have been widely used to investigate the gradient in metallicity with radius in the Galactic disc (e.g., Twarog et al. 2003; Friel et al. 2002, 2010; Magrini et al. 2009; Jacobson et al. 2011a,b). Friel (1995) reviewed the firsts investigations in this field. Since then, a great effort have been performed to obtain both homogeneous (e.g., Friel et al. 2002, 2010; Sestito et al. 2008) and/or larger samples (e.g., Twarog et al. 1997; Jacobson et al. 2011a,b). All these investigations agree on the fact that the iron content decreases with increasing radius (e.g., Friel et al. 2002). This behaviour has been generally considered linear with a slope between  − 0.05 and  − 0.09 dex kpc-1, depending on the cluster sample used. Similar trends were obtained for other different tracers of the disc (e.g., Andrievsky et al. 2004; Lemasle et al. 2008). Most of these works were limited to the inner Rgc  ≃  15 kpc. However, investigations based on samples containing clusters at larger distances (e.g., Twarog et al. 1997; Yong et al. 2005; Sestito et al. 2008) found that the [Fe/H] ratio decreases as a function of increasing radius to Rgc  ≃  12.5 kpc and appears to flatten from there outwards.

thumbnail Fig. 8

Trends of [Fe/H] (top panel) and [α/Fe] (bottom panel) with galactocentric radius. Grey dots are OC compiled in Table 12, while black dots are the ones analysed here and in Paper I. A global linear fit is drawn in both panels (long-dashed line). Two separate linear fits of OC inside and outside 12.5 kpc are also shown (solid lines).

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The variation in [Fe/H] with Rgc in our compilation has been plotted in the top panel of Fig. 8. The whole sample is well fitted by a line with a slope of  − 0.046  ±  0.005 dex kpc-1 (long-dashed line), in concordance with the result obtained in Paper I from a  ≃20% smaller sample (−0.05  ±  0.01 dex kpc-1) and in other investigations in the literature (e.g.,  −0.06  ±  0.02 dex kpc-1; Friel et al. 2002). The sample used here contains more clusters with distances larger than Rgc  ≥  12 kpc. This allows us to investigate the discontinuity observed by some authors at Rgc  ≃  12 − 13 kpc. At first sight, no clear discontinuity in slope appears, partly because of the large range of [Fe/H] at this radius ( ≃ 0.5 dex) and partly as a possible consequence of the heterogeneity of our sample. However, when we fit separately clusters inwards and outwards of 12.5 kpc, we find two significantly different slopes: the metallicity in the inner disc decreases with a slope of  − 0.07  ±  0.01 dex kpc-1, while in the outer disc the slope is  − 0.01  ±  0.01 dex kpc-1. The obtained slopes change within the uncertainties if the cut radius varies between 11.5 and 13.5 kpc. This is also in very good agreement with the recent results by Andreuzzi et al. (2011), who find  − 0.07 dex kpc-1 in the inner 12 kpc. This bimodal behaviour can be explained by a different chemical enrichment and star formation in the inner and outer disc; (e.g., Chiappini et al. 2001; Magrini et al. 2009) however, a sharp discontinuity between the inner and outer disc is not expected theoretically.

The ratio [α/Fe] reflects the relative contributions of type Ia and II supernovae: chemical evolution models predict an increase of this ratio with Rgc (e.g., Chiappini et al. 2001; Magrini et al. 2009). This tendency was indeed observed in OC by, e.g., Yong et al. (2005), Magrini et al. (2009), and in Paper I. The bottom panel of Fig. 8 shows the variation in [α/Fe] with Rgc for our compilation: a weak increase in α-element abundances with radius is apparent. However, the slope is still compatible with a flat distribution at the 1σ level, as in Paper I, especially if the two outermost clusters are removed. The discontinuity observed for [Fe/H] is not evident at all in [α/Fe].

An accretion of a satellite into the outer disc could also explain the trend observed (e.g., Chiappini et al. 2001; Yong et al. 2005). In this case, we would expect to find some inhomogeneities corresponding to the trajectory of the merger. Carraro & Bensby (2009) indeed found evidence that two OC, Berkeley 29 and Saurer 1, are related to the Sagittarius dwarf galaxy. Our compiled sample unfortunately do not allow us to investigate this question in depth.

thumbnail Fig. 9

Gradient in [Fe/H] as a function of Rgc in four different age bins (labeled in top-right corner). A linear fit is performed for the OC within a radius of Rgc = 12.5 kpc, and the slope indicated on each panel. A flatter and roughly constant slope is found outside a radius of Rgc = 12.5 kpc.

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6.2. Time evolution of the radial gradient

Chemical evolution models of the Galactic disc predict a variation in the metallicity gradient with time, but they disagree about the direction of this gradient variation (see Maciel et al. 2007, for a recent review), some predicting a steepening and some a flattening of the gradient with time. Studies based on metallicities derived from low-resolution spectroscopy found that old OC (≳1 Gyr) followed a steeper radial gradient,  ~−0.08 dex kpc-1, than the younger ones,  ~−0.02 dex kpc-1 (Friel et al. 2002; Chen et al. 2003). Only recently have chemical abundances been derived from high-resolution spectroscopy for a sufficient number of OC to significantly investigate the variation in the radial gradient with time. As for studies based on low-resolution spectra, they agree that the gradient was steeper in the past and has flattened with time (Magrini et al. 2009; Andreuzzi et al. 2011). For example, on the basis of a sample of  ~70 OC Andreuzzi et al. (2011) found that all objects younger than 4 Gyr display a similar gradient with a slope  −0.07 dex kpc-1 in the inner 12 kpc, while the one for older objects is steeper,  −0.15 dex kpc-1.

Other tracers have been used to study the time variation in radial gradients. Studies based on planetary nebulae found more puzzling results: while Maciel et al. (2003) found a flattening of the gradient with time, as generally observed for OC, Stanghellini & Haywood (2010) found that the gradient steepens with time. At the moment, there is no explanation of this contradictory result. Comparisons among the slopes of the radial gradients described by populations of different ages also show that the gradient has flattened out in the past few Gyr (see Maciel & Costa 2009, for a recent review).

To investigate the behaviour of the radial gradient in our compiled sample of high-resolution abundances, we plotted in Fig. 9 the gradient in [Fe/H] as a function of Rgc in four different age bins. We obtained a linear fit in each age bin for the inner 12.5 kpc, and for the outer range we simply used the same fit as in Fig. 8, owing to the paucity of OC after age binning in this region. We found that the slope of the [Fe/H] gradient increases as we go back in time from  − 0.02  ±  0.01 dex kpc-1 for objects younger than 0.1 Gyr to  −0.10  ±  0.01 dex kpc-1 for clusters older than 2.5 Gyr.

thumbnail Fig. 10

Trends in [Fe/H] with |z| in four radial annuli as indicated on the top-right corner of each panel, moving outwards from the top to the bottom panel. Symbols are the same as in Fig. 8. As a reference, we plotted dashed lines in each panel, representing the median metallicity of clusters in each radial annulus.

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6.3. Trends with the disc scale-height

Another interesting trend that could be investigated is the behaviour of [Fe/H] with the vertical scale-height of the disc z, i.e., the vertical [Fe/H] gradient. Although the formation of the thick discs remains an open question, the existence of vertical gradients can help us to discriminate among the mechanisms proposed to their formation. No vertical chemical gradients are expected in thick discs formed by heating caused by accretion events or major mergers. In contrast, vertical gradients may exist in discs thickened by gradual heating of the thin disc or before the gas has settled to form a thin disc (see Mould 2005, for a review). Up to now, there is no conclusive agreement about the existence of a vertical metallicity gradient in the Galactic disk. The existence of a vertical gradient for field stars have been claimed by several authors, although they cover only about 1 kpc above and below the disc plane (Bartašiūtė et al. 2003; Marsakov & Borkova 2005, 2006; Soubiran et al. 2008). Studies covering large ranges of |z| do not find any evidence of a vertical gradient (Gilmore et al. 1995; Soubiran & Girard 2005; Navarro et al. 2011) among the field populations. Studies using OCs have found a vertical gradient of  ~ −0.3 dex kpc-1(Piatti et al. 1995; Carraro et al. 1998; Chen et al. 2003), although, these studies do not distinguish the effects of the radial gradient, which can mask any vertical trend. This effects were taken into account by Jacobson et al. (2011a) who found no evidence of a vertical gradient.

To investigate the presence of trends with z in our compilation, we firstly had to remove the contribution of the radial metallicity gradient. We plotted in Fig. 10 the variation in [Fe/H] with |z| in four different annuli of Rgc. We note that OC with high |z| are preferentially located at large Rgc; this is not unexpected because the disc thickens in its external regions. Moreover, an intrinsic bias caused by obscuration in the plane appears: clusters at large Galactocentric radii are found and observed preferentially higher above the plane. This could explain why the two outermost annuli studied uncover a possible weak decrease in [Fe/H] as z increases. This trend is however still compatible with no gradient at the 1σ level and, once again, larger samples of homogeneous data are necessary to investigate this result in detail.

thumbnail Fig. 11

The evolution of [Fe/H] with age in the same four radial annuli as in Fig. 10. Again, dashed lines representing the median metallicity of clusters in each radial annulus have been plotted as reference.

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6.4. Is there an age-metallicity relation for open clusters?

Another important prediction of the chemical evolution models is the existence of an age-metallicity relation for disc populations. It is still unclear whether or not the field disc stars follow an age-metallicity relation. Some works find it (e.g., Reddy et al. 2003; Bensby et al. 2004; Reid et al. 2007), but others do not (e.g., Feltzing et al. 2001; Nordström et al. 2004; Karataş et al. 2005). Again, no clear trend of chemical abundances with age has been clearly observed in the case of Galactic OC (e.g., Friel 1995). Although Friel et al. (2010) notices a trend of [Al/Fe] and [O/Fe] ratios with age, again larger and homogeneous samples are necessary to confirm this result. If an age-metallicity relation is confirmed for the field population but not for OC, this would imply that they might have followed a different chemical evolution (Yong et al. 2005).

The evolution of the radial gradient as a function of time, described above, indicates that the chemical enrichment of OC is modulated by their location in the Galaxy and not by the moment at which they formed. To investigate whether an age-metallicity relationship exits at a given Rgc, we plotted in Fig. 11 the evolution of [Fe/H] with age in four different radial annuli. There is no clear trend in any of the studied annuli, although not all of them contain clusters covering the same age range. Only in the outermost annulus is a weak trend observed, although it is still not very significant. Again, we conclude that a larger sample of homogeneous data are necessary to investigate this point in depth.

7. Summary and conclusions

We have enlarged our sample of homogeneous high-resolution abundance measurements from the five clusters of Paper I to a total of nine, analysing here spectra of red clump giants in the Hyades, Praesepe, NGC 752, and Be 32. Our main results can be summarized as follows:

  • We provide the first high-resolution based abundance ratios(other than [Fe/H], see Sestitoet al. 2008) forNGC 752, which turned out to be mostly of solarcomposition.

  • We have presented the abundance ratios of Praesepe red clump giants, which appear to solve a puzzling dichotomy of literature determinations for stars of different evolutionary stages.

  • We have found that our abundance ratios for the Hyades and Berkeley 32 are in good agreement with other literature determinations.

  • We have confirmed the absence of light elements (anti-)correlations in the OC studied so far.

We have updated our compilation of previous literature data for 57 clusters of Paper I to a total of 89 clusters presented here. With this updated compilation and our homogeneous measurements in hand, we have investigated Galactic trends in [Fe/H] (and [α/Fe]) with age, Galactocentric radius, and height above the Galactic plane. Our findings are in substantial agreement with other similar investigations, where the abundance gradient appears to indeed flatten out outside Rgc  ≃  12.5 kpc, and the inner disc slope appears to flatten for younger ages as well, although the age bins are not too well-sampled. At the same time, [α/Fe] shows a weak increase with Rgc. No significant gradients are observed with |z| or age, except for a weak tendency of [Fe/H] to decrease with increasing |z| and decrease with age in the outermost disc annulus studied. None of our measured weak trends have any significance above 1σ. Larger samples of homogeneous data are still necessary to investigate the existence of any dependence on age and |z| in the Galactic disc.


2

http://www.ipac.caltech.edu/2mass. 2MASS (Two Micron All Sky Survey) 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.

3

Image Reduction and Analysis Facility, IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

5

However, see the discussion by Mucciarelli (2011) about the pros and cons of the Magain (1984) correction, which depends heavily on data quality and line selection effects.

7

We decided not to use the Magain (1984) effect, because we prefer to have internally consistent abundances from each line, and because of the additional effects described by Mucciarelli (2011).

Acknowledgments

We acknowledge the anonymous referee for helping us to improve this paper. R.C. acknowledge the support from the Spanish Ministry of Science and Technology (Plan Nacional de Investigación Científica, Desarrollo, e Investigación Tecnológica, AYA2004-06343 and AYA2007-3E3507). R.C. also acknowledges the funds by the Spanish Ministry of Science and Innovation under the Juan de la Cierva and MEC/Fullbrigth fellowships, and by the Centro de Investigaciones de Astronomía (Venezuela) under its postdoctoral fellowship programe.

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Online material

Table 12

Literature sources of high-resolution (R  ≥  1500) [Fe/H] ratios of open clusters together with the resolution, signal-to-noise ratios, number of stars, and method used in each of them.

All Tables

Table 1

Observing logs and programme star properties.

Table 2

Heliocentric radial velocity measurements and 1σ errors (Vr ± δVr)here for each programme star.

Table 3

Adopted cluster parameters.

Table 4

Stellar atmosphere parameters for the programme stars (see text).

Table 5

Equivalent widths and atomic data of the programme stars.

Table 6

Abundance ratios for single cluster stars, with their internal and external (last column) uncertainties.

Table 7

Average cluster abundances, obtained as the weighted average of the single stars abundances in each of them.

Table 8

High-resolution average Be 32 abundances.

Table 9

Abundance comparison of individual Hyades stars (see text).

Table 10

High-resolution average Hyades abundances from giants.

Table 11

High-resolution average Praesepe (NGC 2632) abundances.

Table 12

Literature sources of high-resolution (R  ≥  1500) [Fe/H] ratios of open clusters together with the resolution, signal-to-noise ratios, number of stars, and method used in each of them.

All Figures

thumbnail Fig. 1

Location of target stars (large black dots with star ID labels) in the color–magnitude diagrams of their respective parent clusters (small grey dots).

Open with DEXTER
In the text
thumbnail Fig. 2

Comparison of our EW measurements with those by Bragaglia et al. (2008) for star 0456 in Be 32, and by Boyarchuk et al. (2000) for three Hyades giants. Dotted lines mark perfect agreement (zero difference), while dashed lines are linear fits to the data.

Open with DEXTER
In the text
thumbnail Fig. 3

Comparison between our iron-peak abundance ratios (large black dots), those of Paper I (large black open circles), high-resolution measurements listed in Table 12 (large dark grey dots), field stars belonging to the thin disc (light grey dots, Reddy et al. 2003), and to the thick disc (tiny light grey dots, Reddy et al. 2006). Errorbars in our measurements are the quadratic sum of all uncertainties discussed in Sects. 3.2 and 3.3.

Open with DEXTER
In the text
thumbnail Fig. 4

Comparison of our α-elements ratios with literature values. Symbols are the same as in Fig. 3.

Open with DEXTER
In the text
thumbnail Fig. 5

Comparison of our s-process elements ratios with the literature ones. Symbols are the same as in Fig. 3, except for the black star-like symbols in the top [Ba/Fe] panel, which represent the revision of Ba abundances with spectral synthesis performed by D’Orazi et al. (2009).

Open with DEXTER
In the text
thumbnail Fig. 6

Comparison between our [Na/Fe] and [Al/Fe] ratios and the literature ones. Symbols are the same as in Fig. 3.

Open with DEXTER
In the text
thumbnail Fig. 7

A search for (anti)-correlations of Al, Mg, Na, and O among our target stars. The four panels show different planes of abundance ratios, where stars belonging to each cluster are marked with different symbols. Dotted lines show solar values, solid lines show linear regressions and the typical uncertainty (~0.1 dex) is marked at the lower right corner of each panel.

Open with DEXTER
In the text
thumbnail Fig. 8

Trends of [Fe/H] (top panel) and [α/Fe] (bottom panel) with galactocentric radius. Grey dots are OC compiled in Table 12, while black dots are the ones analysed here and in Paper I. A global linear fit is drawn in both panels (long-dashed line). Two separate linear fits of OC inside and outside 12.5 kpc are also shown (solid lines).

Open with DEXTER
In the text
thumbnail Fig. 9

Gradient in [Fe/H] as a function of Rgc in four different age bins (labeled in top-right corner). A linear fit is performed for the OC within a radius of Rgc = 12.5 kpc, and the slope indicated on each panel. A flatter and roughly constant slope is found outside a radius of Rgc = 12.5 kpc.

Open with DEXTER
In the text
thumbnail Fig. 10

Trends in [Fe/H] with |z| in four radial annuli as indicated on the top-right corner of each panel, moving outwards from the top to the bottom panel. Symbols are the same as in Fig. 8. As a reference, we plotted dashed lines in each panel, representing the median metallicity of clusters in each radial annulus.

Open with DEXTER
In the text
thumbnail Fig. 11

The evolution of [Fe/H] with age in the same four radial annuli as in Fig. 10. Again, dashed lines representing the median metallicity of clusters in each radial annulus have been plotted as reference.

Open with DEXTER
In the text

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