4 Relative abundances

The abundances relative to hydrogen [A/H] and $\sigma$ (the line-to-line scatter) derived for up to 26 neutral and ionized species for the programme stars are listed in Table 4. The abundances of barium are corrected for non-LTE effects by the subtraction of 0.20 dex (see Sect. 4.5 for discussion).

4.1 Carbon and nitrogen

The carbon abundances obtained in our work were compared with carbon abundances determined for dwarf stars in the galactic disk. Gustafsson et al. (1999), using the forbidden [C I] line, performed an abundance analysis of carbon in a sample of 80 late F and early G type dwarfs. Since carbon abundances obtained using the [C I] 8727 Å line and ${\rm C}_2$ molecular lines are usually consistent (cf. Clegg et al. 1981), we expect no systematic shift to be present because of the different abundance indicators used. As is seen from Fig. 4, the ratios of [C/Fe] in our stars lie much below the trend obtained for dwarf stars in the Galactic disk (Gustafsson et al. 1999).

Abundances in the investigated stars suggest that carbon is depleted by about 0.3 dex and nitrogen is enhanced by more than 0.4 dex. These abundance alterations of carbon and nitrogen are larger than those we obtained for the clump stars in the old, solar-metallicity open cluster M 67 (Paper I), but smaller than was found for more metal deficient RHB stars by Gratton et al. (2000a). This brings additional evidence that mixing processes are metallicity dependent. The C/N ratios in the investigated stars are lowered to values in the range 0.7 to 1.2 which is less than predicted by present day stellar evolution calculations. Gratton et al. (2000a) receive even smaller C/N ratios for the two red horizontal branch stars with [Fe/H] about -1.5 dex. The $^{12}{\rm C}/^{13}{\rm C}$ ratios are lowered and lie between values 3 and 7 which indicate extra-mixing processes to be quite strong. Six more metal-deficient RHB stars investigated by Gratton et al. (2000a) show $^{12}{\rm C}/^{13}{\rm C}$ ratios from 6 to 12.

The theoretical standard stellar evolution of the surface carbon isotopic ratios and carbon to nitrogen ratios along the giant branch was homogeneously mapped by Charbonnel (1994) and more recently by Girardi et al. (2000) for stellar masses between 1 and $7~M_{\odot}$ and different metallicities. Our investigated field RHB stars are somewhat metal deficient ( $Z\approx0.008$) and have masses approximately 0.8 to $0.9~M_{\odot}$ (Tautvaisiene 1996b). According to Girardi et al. (2000), the C/N and ¹²C/¹³C ratios in such stars should drop after the first dredge-up episode to values of about 3 and 35, respectively. Charbonnel (1994, extrapolation to ${\approx} 0.85~M_\odot$ in Figs. 2 and 4) predicted similar values after the first dredge-up.

It has long been known that giant stars regularly show much larger evolutionary changes in these abundances than standard models predict, see e.g. Boothroyd & Sackmann (1999) for references. This is the case also for our derived ¹²C/¹³C and C/N ratios. Because of grave differences between model predictions and observations, Charbonnel (1995), Charbonnel et al. (1998) and Boothroyd & Sackmann (1999) performed calculations of models with deep mixing after the first dredge-up. Boothroyd & Sackmann e.g. fitted a one-parameter recipe for "cool bottom processing'' (CBP) after the first dredge-up to the available observations of red-giant abundances. Their CBP results are given for initial stellar masses above  $1.0~M_\odot$. It is difficult to say what were the initial masses of the stars we investigate. It could be that they lost about 0.1-0.3 $M_\odot$ during their evolution on the giant branch (Renzini 1981; Renzini & Fusi Pecci 1988).

The $^{12}{\rm C}/^{13}{\rm C}$ ratios determined for the investigated stars are in quite good agreement with "cool bottom processing'' predictions (Boothroyd & Sackmann 1999) for low mass stars with Z=0.007. The C/N ratios, however, request the initial mass of the stars to be of about $1.8~M_\odot$. The metal-deficient RHB stars investigated by Gratton et al. (2000a) show higher than predicted by CBP $^{12}{\rm C}/^{13}{\rm C}$ ratios but even lower C/N ratios. The low C/N ratios may be an indication that CBP is stronger in such stars than the metallicity scaling of models suggest. However, in view of the sensitivity of C/N ratios to the carbon abundances, we will not claim that the C/N predictions of Sackmann & Boothroyd are wrong, but rather that the C and N abundances should be checked in further studies employing other atomic and molecular features.

 

 

Table 4: Abundances relative to hydrogen [A/H] derived for programme stars. The quoted errors, $\sigma$, are the standard deviations in the mean value due to the line-to-line scatter within the species. Details on error estimates of the ¹²C/¹³C ratios are described in Sect. 3. The number of lines used is indicated by n.

	BD+25$^\circ$2436				BD+25$^\circ$2459				BD+25$^\circ$2502				BD+27$^\circ$2057
Ion	[A/H]	$\sigma$	n		[A/H]	$\sigma$	n		[A/H]	$\sigma$	n		[A/H]	$\sigma$	n
C (C2)	-0.63		1		-0.58		1				1		-0.78		1
N (CN)	-0.01	0.14	154		0.14	0.11	111						-0.27	0.15	102
O I	-0.25		1		-0.24		1		-0.24		1		-0.25		1
Na I	-0.54	0.01	2		-0.28	0.15	2		-0.72	0.02	2		-0.40	0.10	2
Mg I	-0.21	0.04	2		-0.24	0.09	2		-0.32		1		-0.20	0.10	2
Al I	-0.47	0.06	4		-0.34	0.08	4		-0.68	0.05	3		-0.45	0.03	4
Si I	-0.31	0.12	14		-0.18	0.12	13		-0.48	0.10	6		-0.24	0.11	14
Ca I	-0.36	0.12	7		-0.23	0.12	8		-0.53	0.14	4		-0.39	0.15	7
Sc I	-0.40	0.10	4		-0.31	0.11	4		-0.39		1		-0.56	0.15	4
Sc II	-0.33	0.11	10		-0.20	0.08	11		-0.43	0.16	6		-0.46	0.11	10
Ti I	-0.29	0.12	23		-0.11	0.15	21		-0.57	0.16	6		-0.32	0.13	23
Ti II	-0.26		1		-0.05		1						-0.33		1
V I	-0.40	0.10	17		-0.22	0.12	18		-0.70	0.10	6		-0.50	0.15	18
Cr I	-0.38	0.04	7		-0.23	0.13	7						-0.55	0.09	7
Mn I	-0.46	0.08	3		-0.31	0.09	3						-0.70	0.08	2
Fe I	-0.48	0.12	43		-0.35	0.10	42		-0.74	0.06	18		-0.60	0.12	40
Fe II	-0.48	0.10	5		-0.35	0.13	5		-0.74	0.11	2		-0.60	0.12	5
Co I	-0.41	0.13	10		-0.30	0.14	9		-0.58	0.04	2		-0.47	0.14	8
Ni I	-0.42	0.14	22		-0.24	0.13	20		-0.80	0.09	8		-0.56	0.15	21
Y I	-0.37		1		-0.28		1						-0.53		1
Y II	-0.44	0.03	3		-0.29	0.06	3						-0.46	0.11	2
Zr I	-0.50	0.11	3		-0.32	0.13	4		-0.41	0.06	3		-0.67	0.13	4
Ba II	-0.70	0.07	2		-0.34	0.04	2		-0.69	0.05	2		-0.82	0.03	2
La II	-0.62		1		-0.38		1						-0.68		1
Sm II	-0.53		1		-0.05		1						-0.36		1
Eu II	-0.05		1		-0.05		1						-0.15		1
C/N	0.96				0.76								1.23
12C/13C	5	+5/-2			7	+3/-2							5	+2/-2

 

Table 4: continued.

	BD+28$^\circ$2079				BD+29$^\circ$2231				BD+29$^\circ$2294				BD+29$^\circ$2321
Ion	[A/H]	$\sigma$	n		[A/H]	$\sigma$	n		[A/H]	$\sigma$	n		[A/H]	$\sigma$	n
C (C2)					-0.60		1		-0.80		1		-0.68		1
N (CN)					0.15	0.12	162		-0.10	0.14	34		0.00	0.13	93
O I	-0.10		1		-0.22		1		-0.31		1
Na I	-0.51		1		-0.34	0.09	2		-0.28		1		-0.35		1
Mg I	-0.21		1		-0.18	0.02	2		-0.14		1		-0.26		1
Al I	-0.22	0.03	2		-0.32	0.08	4		-0.46	0.06	2		-0.29	0.12	2
Si I	-0.20	0.12	6		-0.20	0.07	14		-0.30	0.08	8		-0.27	0.08	8
Ca I	-0.22	0.13	5		-0.29	0.17	8		-0.39	0.11	5		-0.41	0.14	7
Sc I					-0.26	0.10	4		-0.48	0.07	3		-0.50	0.15	3
Sc II	-0.19	0.07	8		-0.23	0.09	10		-0.32	0.10	9		-0.33	0.08	8
Ti I	-0.12	0.17	16		-0.18	0.12	22		-0.27	0.13	19		-0.43	0.14	20
Ti II	-0.14		1		-0.30		1		-0.25		1		-0.42		1
V I	-0.17	0.16	12		-0.18	0.12	17		-0.43	0.14	12		-0.49	0.13	14
Cr I	-0.27	0.13	7		-0.39	0.15	8		-0.51	0.13	7		-0.51	0.15	7
Mn I	-0.32	0.12	2		-0.32	0.16	3		-0.70	0.05	2		-0.57	0.08	3
Fe I	-0.44	0.06	25		-0.39	0.12	43		-0.54	0.10	24		-0.50	0.08	27
Fe II	-0.44	0.12	3		-0.39	0.07	5		-0.54	0.12	3		-0.50	0.08	3
Co I	-0.35	0.17	4		-0.28	0.16	9		-0.37	0.12	6		-0.41	0.08	7
Ni I	-0.38	0.10	15		-0.33	0.15	22		-0.45	0.11	21		-0.50	0.12	21
Y I					-0.41		1		-0.53		1		-0.57		1
Y II	-0.43	0.14	2		-0.45	0.11	4		-0.64	0.01	2		-0.56	0.13	4
Zr I	-0.42	0.04	2		-0.29	0.02	3		-0.42	0.15	4		-0.53	0.17	3
Ba II	-0.36	0.03	2		-0.54	0.01	2		-0.67		1		-0.60		1
La II					-0.47		1
Sm II	-0.06		1		-0.20		1						-0.25		1
Eu II					-0.05		1		-0.28		1		-0.22		1
C/N					0.71				0.79				0.83
12C/13C					3.5	+3/-1.5			3	+2/-1			4	+4/-1

 

Table 4: continued.

	BD+33$^\circ$2280				BD+34$^\circ$2371				BD+36$^\circ$2303				HD 104783
Ion	[A/H]	$\sigma$	n		[A/H]	$\sigma$	n		[A/H]	$\sigma$	n		[A/H]	$\sigma$	n
C (C2)					-0.45		1		-0.90		1		-0.72		1
N (CN)					0.13	0.13	38		-0.34	0.12	92		-0.05	0.14	109
O I	-0.29		1		-0.13		1						-0.14		1
Na I	-0.33	0.02	2		0.02		1		-0.72		1		-0.56	0.11	2
Mg I	-0.28		1		-0.05		1		-0.36		1		-0.23	0.06	2
Al I	-0.41	0.04	4		-0.10	0.07	2		-0.55	0.05	2		-0.54	0.04	4
Si I	-0.29	0.09	7		-0.06	0.06	7		-0.37	0.12	8		-0.24	0.08	15
Ca I	-0.47	0.19	3		-0.07	0.15	7		-0.55	0.13	7		-0.34	0.13	7
Sc I	-0.34		1		-0.10		1		-0.76	0.03	2		-0.56	0.09	2
Sc II	-0.33	0.07	6		-0.06	0.11	9		-0.48	0.12	9		-0.41	0.09	11
Ti I	-0.39	0.08	6		-0.03	0.13	17		-0.51	0.11	20		-0.23	0.14	24
Ti II					0.07		1		-0.48		1		-0.19		1
V I	-0.44	0.08	6		-0.18	0.09	15		-0.70	0.14	15		-0.44	0.11	16
Cr I	-0.36	0.12	2		-0.22	0.10	6		-0.87	0.13	6		-0.52	0.10	7
Mn I					-0.21	0.18	2		-0.89	0.04	2		-0.63	0.04	2
Fe I	-0.48	0.05	24		-0.18	0.08	29		-0.76	0.08	27		-0.55	0.12	39
Fe II	-0.48	0.07	2		-0.18	0.08	3		-0.76	0.08	3		-0.55	0.10	5
Co I	-0.41	0.16	3		-0.16	0.09	7		-0.67	0.12	7		-0.44	0.12	7
Ni I	-0.42	0.15	9		-0.16	0.09	20		-0.72	0.12	20		-0.48	0.13	22
Y I					-0.24		1						-0.47		1
Y II					-0.13	0.11	3		-0.71	0.09	3		-0.48	0.05	2
Zr I	-0.55	0.07	3		-0.22	0.09	4		-0.70	0.09	3		-0.32	0.06	3
Ba II	-0.47	0.08	2		-0.14	0.09	2		-0.70		1		-0.43	0.09	2
La II	-0.42		1										-0.48		1
Sm II					-0.01		1		-0.55		1		-0.41		1
Eu II					0.00		1		-0.20		1		-0.15		1
C/N					1.05				1.10				0.85
12C/13C					>5				3	+2/-1			>5

4.2 Sodium and aluminium

Sodium and aluminium are among the mixing-sensitive chemical elements. The star-to-star variations of Na, the existence of Na versus N correlations, and Na versus O anticorrelations in globular cluster red giants have revealed the possibility that sodium and aluminium are produced in red giant stars (see Kraft 1994 and Da Costa 1998 for reviews). It is found also that Na variations exist in all clusters, while Al variations are greater in the more metal-poor clusters (cf. Norris & Da Costa 1995; Shetrone 1996, Paper I).

Pilachowski et al. (1996) determined sodium abundances for 60 metal-poor halo subgiants, giants, and horizontal branch stars using high dispersion spectra and concluded that there is an intrinsic difference between halo field giants and globular cluster giants. The bright giants in the field do not show the sodium excesses seen in their globular cluster counterparts. The [Na/Fe] ratios in field stars show a wide scatter (ranging from -0.6to nearly +0.3) with a slight tendency for <[Na/Fe]> to increase with advancing evolutionary stage. In a sample of ten field RHB stars investigated by Tautvaisiene (1997) only two of the more metal rich ([Fe/H]$\,\,>-0.5$) stars showed sodium overabundances of 0.2-0.3 dex.

The stars in our sample show Na and Al abundances which are typical of unevolved stars in the solar vicinity, as determined from the Na I lines $\lambda$ 5682.64 and 6154.23 Å and Al I lines $\lambda$ 6696.03, 6698.66, 7835.31 and 7836.13 Å, see Fig. 5. Gratton et al. (2000a) investigated possible non-LTE effects for the Na I lines, and find the probable corrections not to be larger than about 0.02 dex at the temperatures and gravities of the stars analysed here.

Theoretical explanations for the production of Na and Al have been proposed by Sweigart & Mengel (1979), Langer & Hoffman (1995), Cavallo et al. (1996), Mowlavi (1999), Weiss et al. (2000) and other studies. The nature and extent of the phenomenon is, however, still not well understood.

Prochaska et al. (2000) investigated abundances of Na and Al in 10 thick disk dwarfs and found aluminium to be much more overabundant than sodium. Our sample of thick disk stars does not show such a pattern.

 

Table 4: continued.

	HD 105944
Ion	[A/H]	$\sigma$	n
C (C2)	-0.60		1
N (CN)	0.07	0.13	77
O I	-0.29		1
Na I	-0.24		1
Mg I	-0.17		1
Al I	-0.34	0.02	2
Si I	-0.31	0.11	8
Ca I	-0.19	0.12	6
Sc I	-0.27	0.04	2
Sc II	-0.28	0.11	9
Ti I	-0.28	0.14	19
Ti II	-0.15		1
V I	-0.36	0.11	14
Cr I	-0.41	0.14	7
Mn I	-0.40	0.12	2
Fe I	-0.37	0.08	30
Fe II	-0.37	0.06	3
Co I	-0.38	0.11	8
Ni I	-0.37	0.14	20
Y I
Y II	-0.54	0.06	3
Zr I	-0.29	0.10	2
Ba II	-0.20	0.03	2
La II
Sm II	-0.17		1
Eu II	-0.16		1
C/N	0.85
$^{12}{\rm C}/^{13}{\rm C}$	3.5	+4/-2

Figure 5: [Na/Fe] and [Al/Fe] ratios as a function of iron [Fe/H]. Results for the field RHB stars investigated in the present work are indicated by filled circles, for the Galactic disk stars investigated by Edvardsson et al. (1993) by crosses.

4.3 Oxygen and magnesium

Surface abundances of oxygen and magnesium could be altered in stars only by very deep mixing. E.g., in cluster giants with large aluminium enhancements ($\sim$1.0 dex) produced by very deep mixing, Mg depletions should then be about $\sim$0.2 dex (Langer & Hoffman 1995). Since this is not the case for the investigated stars we will discuss our results for oxygen and magnesium in the context of the thick disk of the Galaxy.

In Figs. 6 and 7, we plot oxygen and magnesium abundance ratios and compare them with the modeled ratios describing the mean trend of the Galactic thin disk (Pagel & Tautvaisiene 1995). Other results obtained for the thick disk stars in recent studies are displayed as well. Prochaska et al. (2000) analysed a sample of 10 thick disk stars with the HIRES spectrograph on the 10 m Keck I telescope. Unfortunately, the forbidden [O I] $\lambda 6300$ Å line fell in the inter-order gap and the less trustworthy O I triplet lines at 7775 Å had to be used in their analysis. We adopt for the figures the results for 4 thick disk stars from the work by Gratton et al. (2000b). In the same paper a sample of thick disk candidates was selected from the work by Edvardsson et al. (1993). Stars which have [O/H]$\,>-0.5$, [Fe/O]$\,<-0.25$ and $-0.5<\,$[Mg/H]$\,<0$, [Fe/Mg]$\,<-0.25$and appropriate dynamical parameters were attributed to the thick disk. While plotted, the data make quite a cloud lying above the semiempirical trends modeled for the thin disk of the Galaxy by Pagel & Tautvaisiene (1995), but this can hardly be used to draw any conclusions about the location in terms of metallicity of the transition between the halo and thick disk populations. The high accuracy results for magnesium determined by Fuhrmann (1998) lie at the edge of the distribution. This may be taken as an indication that the transition between the halo phase and the thick disk phase took place around [Fe/H] $\approx -0.6$ to -0.5. Our oxygen and magnesium to iron ratios tend to indicate the onset of supernova of Type Ia (SN Ia) at about [Fe/H] =-0.7 to -0.6. We suggest that a model for the halo and thick disk may look much like the model of Pagel & Tautvaisiene (1995), with the difference that the halo phase continued all the way up to [Fe/H] $\approx -0.6$ dex.

Figure 6: [O/Fe] and [Mg/Fe] ratios as a function of iron [Fe/H] for the thick disk stars analysed in recent studies: filled circles - the present work; triangles - Prochaska et al. (2000); rhombs - Gratton et al. (2000b); crosses - Edvardsson's et al. (1993) dwarfs with $R_{\rm m}\le 7$kpc, reanalysed and selected to be the thick disk stars by Gratton et al. 2000b; squares - Fuhrmann (1998). The solid lines show the model of the Galactic thin disk (Pagel & Tautvaisiene 1995).

Figure 7: Run of [Fe/O] vs. [O/H] and [Fe/Mg] vs. [Mg/H] ratios for the stars of Fig. 6.

4.4 Silicon, calcium and titanium

The $\alpha$-elements silicon, calcium and titanium may also bring information on the thick disk of the Galaxy. A large number of spectral lines with accurate gf-values are available for the analysis which should provide for good abundance precision. Being produced both in Type II and Ia supernova, Si, Ti and Ca may be expected to show smaller overabundances than O and Mg. As is seen from Fig. 8, abundance ratios of these elements to iron may also exhibit slight overabundances with respect to the mean trend of the thin disk.

Figure 8: [Si/Fe], [Ca/Fe] and [Ti/Fe] ratios as a function of iron [Fe/H] for the thick disk stars analysed in recent studies. The meaning of symbols as in Fig. 6.

4.5 s- and r-process elements

As already mentioned, the barium abundances in our study are corrected for non-LTE effects by the subtraction of 0.20 dex. Two quite similar Ba II lines $\lambda 6141$ and 6496 Å were used for the analysis. According to Mashonkina et al. (1999) and Mashonkina & Gehren (2000), the non-LTE correction for the Ba II line $\lambda 6496$ is -0.2 dex on average in the metallicity range $-1 <{\rm [Fe/H]}<0.1$. Non-LTE effects for the line $\lambda 6141$ were not studied well enough, since this line is too saturated in the solar spectrum to provide an accurate correction. Theoretical non-LTE calculations show that non-LTE effects for this line are not smaller than for $\lambda 6496$, only the weak line $\lambda 5853$ Å is quite insensitive. In our study, both $\lambda 6141$ and 6496 Å gave approximately the same barium abundances, so we decided to apply the same correction to both. In the work by Prochaska et al. (2000) three Ba II lines $\lambda 5853$, 6141 and 6496 were used, and a typical correction of 0.17 dex was applied.

Figure 9: Abundance ratios of the s-process dominated (Y, Zr, Ba and La) and r-process dominated (Sm and Eu) elements to iron as a function of iron [Fe/H] for the thick disk stars analysed in recent studies. The meaning of symbols as in Fig. 6, open circles represent results by Mashonkina & Gehren (2000). The solid line shows the model of the Galactic thin disk (Pagel & Tautvaisiene 1997).

Abundance ratios of s- and r-process-dominated (in the Solar system, Burris et al. 2000) elements to iron as a function of iron [Fe/H] for the thick disk stars analysed in the recent studies are presented in Fig. 9. For a comparison, the modeled abundance trends of the Galactic thin disk by Pagel & Tautvaisiene (1997) are shown. As is the case for oxygen and the $\alpha$ elements, these elements fit the models for the thin disk reasonably well if we shift the onset of SN Ia from [Fe/H] =-1.1 to -0.6 dex. Since europium is an almost pure r-process element and supposedly produced with oxygen and magnesium in stars exploding as core-collapse supernovae, the thick-disk Eu abundance trend differ quite dramatically from the thin-disk one and may be very useful for population studies. [Eu/Fe] ratios obtained in our sample of thick disk stars and in ten more stars analysed by Prochaska et al. (2000) and Mashonkina & Gehren (2000) bring quite a clear indication that the thick disk population is chemically discrete from the thin disk.