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The r-Process Enriched Low Metallicity Giant HD 115444

New high resolution, very high signal-to-noise spectra of ultra-metal-poor (UMP) giant stars HD 115444 and HD 122563 have been gathered with the High-Resolution Echelle Spectrometer of the McDonald Observatory 2.7m Telescope. With these spectra, line identification and model atmosphere analyses have been conducted, emphasizing the neutron-capture elements. Twenty elements with Z>30 have been identified in the spectrum of HD 115444. This star is known to have overabundances of the neutron-capture elements, but it has lacked a detailed analysis necessary to compare with nucleosynthesis predictions. The new study features a line-by-line differential abundance comparison of HD 115444 with the bright, well-studied halo giant HD 122563. For HD 115444, the overall metallicity is [Fe/H]~ -3.0. The abundances of the light and iron-peak elements generally show the same pattern as other UMP stars (e.g. overdeficiencies of manganese and chromium, overabundances of cobalt), but the differential analysis indicates several nucleosynthesis signatures that are unique to each star.


Introduction
The secrets of early Galactic nucleosynthesis are slowly being revealed from the chemical compositions of metal-poor stars. Advances in analytical techniques have been applied to increasing samples of fainter (thus more distant) stars, yielding more sharply defined abundance trends with metallicity for several major element groups. The general abundance patterns (and their exceptions) in stars of the metallicity range 0.0 ≥ [Fe/H] ≥ -2.5 4 are now reasonably well established for most spectroscopically accessible elements with atomic numbers Z ≤ 30. A good recent summary of these trends has been given by Cayrel (1996).
Elements with Z > 30, comprising more than half of the Periodic Table, are synthesized almost exclusively in neutron-capture fusion reactions. Abundance trends with metallicity are on firm ground for relatively few of these neutron-capture elements, primarily due to their small abundances and lack of strong transitions available to ground-based spectroscopy for most of these elements. Nevertheless some general statements may be made from the extant observations. First, the "bulk" levels of neutron-capture elements vary greatly with respect to the Fe-peak elements in the lowest metallicity stars ([Fe/H] ∼ < -2.5), hereafter called ultra-metal-poor (UMP) stars. The neutron-capture-to-Fe ratios can scatter by more than a factor of 100 from star to star at a given metallicity (Gilroy et al. 1988, Burris et al. 1999 Burris et al. 1999). Europium is synthesized most easily in conditions of extremely high neutron fluxes (during the so-called r[apid]-process) while barium is most easily made in very low neutron-flux environments (during the s[low]-process). The decline in [Ba/Eu] is thus taken as evidence for the dominance of a Type II supernova-induced r-process in producing neutron-capture elements early in the Galaxy's nucleosynthesis history. Third, the 4 We adopt the usual spectroscopic notations that [A/B] ≡ log 10 (N A /N B ) star -log 10 (N A /N B ) ⊙ , and that log ε(A) ≡ log 10 (N A /N H ) + 12.0, for elements A and B. Also, metallicity will be assumed here to be equivalent to the stellar [Fe/H] value.
Confidence in the r-process origin of the heavier neutron-capture elements in UMP stars does not rest solely with the many determinations of [Ba/Eu] ratios, which are still subject to reasonably large observational uncertainties. Instead, the most convincing evidence has come from the detailed abundance distributions of neutron-capture elements in those UMP stars that are bright enough to permit very high resolution, high signal-to-noise spectroscopy (so that very weak neutron-capture element transitions may be detected), or those that have anomalously large [n-capture/Fe] ratios (so that the neutron-capture transitions are detected in high contrast with features of other elements). Such detailed abundance distributions (at least 10 neutron-capture elements) exist for only a few UMP stars. For example, Sneden & Parthasarathy (1983) published a detailed neutron-capture element study of the bright UMP giant HD 122563. These elements are actually underabundant with respect to Fe in HD 122563 (e.g., [Eu/Fe] ≃ -0.4). Another UMP giant star, CS 22892-052, originally identified in the objective prism survey of the halo by Beers, Preston, & Shectman (1985, 1992, is a much fainter star at V = 13.1. But it has been the target of several recent abundance studies (Sneden et al. 1994Norris, Ryan, & Beers 1997a) because it has the largest overabundances of neutron-capture elements (e.g., [Eu/Fe] ≃ +1.7) of any UMP star outside of the CH-star domain. The common link between the neutron-capture abundance mixes of HD 122563, CS 22892-052, and many other UMP stars studied by Gilroy et al. (1988) and McWilliam et al. (1995) is the r-process signature among the heavy elements.
The relative abundance patterns are remarkably similar, and so far are indistinguishable from the abundance pattern of the r-process-only parts of solar-system material (Käppeler, Beer, & Wisshak 1989).
An important aspect of neutron-capture element studies in UMP stars is nucleocosmochronometry. This is the determination of the ages of halo stars (commonly assumed to be nearly equal to the galactic age) from abundances of elements that radioactively -6decay with long half-lives. The technique was pioneered by Butcher (1987) and has been applied in practice mainly to the spectroscopically accessible element thorium. This element is synthesized exclusively in the r-process and decays with τ 1 / 2 = 14.0 Gyr. But the single strong Th ii feature routinely available to ground-based spectroscopy is part of a complex atomic and molecular transition blend at λ4019Å, as illustrated in spectra published by Morell, Källander, & Butcher (1992), and François, Spite, & Spite (1993). Thorium  gap between thorium and the other elements. A thorium decay age is computed from the difference between the presently observed [Th/n-capture] ratio and the estimated value this ratio had immediately after supernova synthesis. Thus in order to solidify observational aspects of thorium cosmochronometry, many abundances throughout the Z = 60-90 range must be determined in many stars. This means that UMP stars with neutron-capture element overabundances must be sought out and subjected to extensive analyses.
The star HD 115444 was first noted in Bond's (1980) objective prism search for low metallicity halo giants. But its unusual spectrum was discovered by Griffin et al. (1982, hereafter GGGV82), who obtained moderately high resolution photographic spectra of HD 115444 after it failed to yield a sensible measurement with Griffin's (1967) photoelectric radial velocity spectrometer.
Their analysis of HD 115444, performed differentially with respect to HD 122563, concluded that HD 115444 probably was more metal-poor than HD 122563 ([Fe/H] ≃ -3.0 versus -2.7). They also found all six neutron-capture elements detected in both stars to be more abundant in HD 115444 than in HD 122563. However, the heaviest neutron-capture elements Ba, La, and Eu were far more overabundant than the lighter elements Sr, Y, and Zr. GGGV82 labeled HD 115444 as "a barium star of extreme Population II", but noted that the r-process element Eu seemed nearly as overabundant as the s-process element Ba. Gilroy et al. (1988) included HD 115444 in their neutron-capture abundance survey of Bond (1980) giants. Their results, determined from an -7echelle/CCD spectrum, did not indicate an enhancement of barium. Instead, they derived a large europium abundance leading to [Ba/Eu] ≃ -1.0, indicative of an r-process origin for the heavier neutron-capture elements.
Recently, Sneden et al. (1998a) obtained HST GHRS data in three UV spectral regions of HD 115444, and they detected (for the first time in any UMP star) some transitions of the very heavy neutron-capture elements osmium, platinum and lead. Unfortunately, their interpretation of the abundances of these elements was somewhat compromised by the uncertain (possibly conflicting) abundances of other neutron-capture elements from the GGGV82 and Gilroy et al. (1988) studies. In this paper our goals are: a) to clear up questions on the status of HD 115444 remaining from the earlier ground-based studies; b) to expand the number of detected neutron-capture elements in this unusual star; c) to connect in a consistent way its ground-based and HST-based abundances; and d) to make an age estimate for the HD 115444 material from its thorium abundance. We present a comprehensive abundance study of HD 115444 from newly acquired high resolution, high signal-to-noise spectra spanning the near-UV to red spectral regions.
Following the lead of GGGV82, we have analyzed HD 115444 with respect to HD 122563, because the atmospheric parameters of these two halo UMP red giants are so similar. In succeeding sections we describe the data acquisition and reduction, discuss the differential model atmosphere calculations and abundance derivation of these stars, and then concentrate on implications of the neutron-capture element abundances in HD 115444.

Observations, Reductions, and Equivalent Widths
In 1997 we obtained high resolution, high signal-to-noise spectra of HD 115444 and HD 122563 with the McDonald Observatory 2.7m H. J. Smith telescope and the "2d-coudé" cross-dispersed echelle spectrograph (Tull et al. 1995). Basic data for these two stars are given in Table 1. We also obtained companion spectra of a tungsten lamp to use in flat-fielding the raw data frames, a Th-Ar lamp for eventual wavelength calibration of the extracted spectra, and a rapidly rotating hot star for telluric absorption line cancelation. The spectrograph was configured such that with -8an entrance aperture width of 1.2 arcsec in the dispersion dimension, the 2-pixel spectral resolving power at the Tektronix 2048×2048 CCD was R ≡ ∆λ/λ ≃ 60,000. The spectra spanned the wavelength range 3600Å ≤ λ ≤ 10000Å. Blueward of 5900Å the spectral coverage was complete, but redward of this wavelength the free spectral range of the echelle orders was too large to be contained on the CCD, thus producing wavelength gaps in the recorded spectra. The S/N ratios of the spectra were of course wavelength-dependent, such that at λ6000Å, S/N ≃ 500; at λ5000Å, S/N ≃ 400; and at λ4000Å, S/N ≃ 200.
The raw data frames were processed in standard ways using IRAF echelle reduction routines.
After extraction of wavelength calibrated echelle orders, cancelation of telluric spectral features in some of the orders was accomplished with specialized software (Fitzpatrick & Sneden 1987).
Sample spectra of HD 115444 and HD 122563 are shown in Figure 1. In general, absorption features in HD 122563 are somewhat stronger than in HD 115444. This makes sense, given previous claims that HD 122563 is more metal-rich and has a lower effective temperature than HD 115444. However, as seen in the upper panel, the transitions of the "heavy" (Z ≥ 56) neutron-capture elements La and Eu are clearly much stronger in HD 115444. On the other hand, the transitions displayed in the lower panel shows no corresponding enhancements of the "lighter" neutron-capture elements Zr and Sr in HD 115444.
Equivalent widths (EWs) were measured in order to derive model atmosphere parameters and to determine abundances for those elements possessing many unblended lines in our spectra. We measured the EWs with the IRAF task splot, using a Gaussian assumption for the line profiles.
Two continuum points were marked and a single best-fit line profile was displayed and examined for obvious blends. Listed in Table 2 are the measured EWs of HD 115444 and HD 122563 together with the excitation potentials (EPs) and transition probabilities (log gf's) of the lines.
For most lines we adopted the log gf's employed by McWilliam et al. (1995) and Sneden et al. (1996); see those papers for the sources of these data. We also supplemented some log gf's from Fuhr & Wiese (1996). Laboratory gf values have been employed throughout this study; none have been derived from a solar spectral analysis. For those lines that were analyzed using synthetic -9spectra, the entry syn is used in Table 2 instead of an EW value.
The standard deviations are dominated by a couple of very discrepant lines. Gratton & Sneden (1990) also have a small number of lines in common with our study, but only for HD 122563. Their values agree well with our measurements: < δEW> = +1.3 ± 1.4 mÅ, (σ = 4.65 mÅ, 11 lines).
Since our new spectra are of higher S/N and resolution than (most of) the compared studies, the encountered differences were not unexpected. On the whole our EWs agree reasonably well with those in the literature.

Derivation of Fundamental Model Parameters
We first derived stellar model atmospheres and abundances for HD 115444 and HD 122563 in two independent analyses. We then followed the spirit of the GGGV82 study by using HD 122563 as a template star to derive more accurate differential model atmosphere parameters for HD 115444. Results of the independent analyses of each star will be labeled "absolute", while the term "relative" will be reserved for results of the differential analysis. properties. MARCS is a modern version of the code described by Gustafsson et al. (1975) and Edvardsson et al. (1993), and it is subject to ongoing refinements. The code, together with the sources of its opacity data, will be described elsewhere. For the current MARCS implementation, modern data for metallic and molecular lines and continua were compiled from many different sources. The mixing-length approximation was used for convection, with parameters α = ℓ/H p = 1.5, ν = 8 and y = 0.076 (Henyey et al. 1965) for our models. An enhancement of 0.4 dex of the α-process elements (from oxygen through calcium) was also assumed.
Final choice among possible trial model atmospheres was made by using them and the measured EWs as input for the current version of the LTE line analysis program MOOG (Sneden 1973). The model parameters were determined as follows. Effective temperature: T eff was iteratively adjusted until the trend of Fe i abundance with excitation potential was minimized.
Microturbulence: v t mostly affects the abundances derived from strong (saturated) lines. Therefore a plot of abundance vs. EW will show a slope if the microturbulence is in error; v t was iteratively determined by minimizing such slopes. Gravity: To find the appropriate log g we required that the abundance determined from ionized species be equal to the abundance derived from neutral species (Fe and Ti were used To illustrate the T eff and v t derivations, Figure 2 shows abundances of Fe i lines derived with the final model of HD 115444 plotted against excitation potential (EP) and the logarithm of the reduced width (log EW/λ). Trends of abundances with EP and log EW/λ have been minimized.
-11 - The mean Fe i abundances agree to the line-to-line scatter limit (σ ≃ 0.1 dex) for lines ranging over nearly 5 eV in EP and over two orders of magnitude in EW.
In order to investigate internal abundance uncertainties we computed the dependence of The abundance changes with these parameter excursions are given in Table 3. Some trends are clear, such as the inevitable rise in all abundances as T eff increases. A metallicity change hardly alters the abundances. Gravity changes affect neutral and ionized species abundances in opposite directions. Larger microturbulence values will decrease the abundances, but this parameter affects just the strong (saturated) lines, and thus is not a large effect for most species in these generally weak-lined stars. Exceptions to this statement are those species exhibiting only saturated lines (e.g., Sr ii, Ba ii). Even here, the entries of Table 3 constitute a worst-case set of internal uncertainties. Many abundance ratios (i.e., [X/Y] for elements X and Y), will be substantially smaller, especially when comparing abundances from the same ionization states.
Furthermore, derivation of T eff and log g through excitation/ionization equilibria considerations inevitably couples them: higher values of T eff also force higher values of log g, which partially cancels any abundance changes.
The main sources of line-to-line abundance scatter within a species for each star are line blending, log gf uncertainties, and EW measurement errors. These errors can be substantially reduced by computing relative abundances between the two stars. We show this in Figure 3 by plotting line-by-line abundances differences ∆ log ε(X) ≡ log ε(X) HD 115444 -log ε(X) HD 122563 with respect to EP and EW, using the abundances generated with the final models for each star. Not only are the slopes negligible (implying well-determined T eff and v t differences), but the line-to-line scatters have shrunk from ≃ 0.1 to ≃ 0.05 dex. In these relative abundances the uncertainties in individual log gf values have been explicitly eliminated. Additionally, errors due to unaccounted-for line blending should affect the derived abundances in the stars in approximately the same way if the blending agents are transitions of elements with Z ≤ 30. Thus we expect a -12further reduction of error in the relative quantities. Some line measurement errors may also be reduced, to the extent that continua have been set and line profiles measured in consistent ways for the two stars. Notice in the lower left-hand panel of Figure 3 that the line-to-line scatter in ∆ log ε(Fe) is largest for the stronger lines (log EW/λ > −4.8). These lines lie on the flat and damping portions of the curve-of-growth, where small EW errors can lead to large abundance errors. But even for these lines the relative Fe abundance is more accurately determined than the absolute abundance of each star.
Choice of model atmosphere code affects the derived stellar parameters and absolute abundance scales. We tested the magnitude of this uncertainty source by deriving atmosphere parameters for the two stars with models interpolated within a recent grid of atmospheres (Kurucz 1992) generated with the ATLAS code. The resulting models are approximately 125 K hotter in T eff , about 0.2 dex higher in log g, but require no obvious changes in microturbulent velocities v t .
The derived abundances average ≃0.06 dex greater with these models. However, model parameter and individual abundance differences between the two stars are virtually unchanged by this switch in model codes.

External Comparisons
Our derived model parameters agree well with previous estimates, especially in the differential sense. There are many more published analyses of HD 122563 in the literature than of HD 115444. -13 -It is more relevant here to discuss those literature studies that treat both of our stars, beginning with effective temperature estimates. Taking differences between the two stars in the sense HD 115444 minus HD 122563, our temperature analyses yielded ∆(T eff ) = 4650 -4500 = +150 K. Three spectroscopic studies have dealt with both of these stars. GGGV82 found ∆(T eff ) = 4800 -4600 = +200 K (and also pointed out that the difference in effective temperatures is more certain than the individual values). Gilroy et al. (1988) Table 1 yield ∆(T eff B−V ) = 4853 -4717 = +136 K. None of the spectroscopic or photometric estimates of T eff differences would claim better than ≃50 K accuracy; thus our newly derived ∆(T eff ) = +150 K between the two stars agrees well with the other values. The absolute T eff scale is less certain, and we note that our temperatures are typically ∼100 K lower than others. This small discrepancy may partly be related to our use of MARCS model atmospheres ( §3.1). But it does not affect any of the basic conclusions of this paper and will not be considered further here.
log g value for HD 115444 is probably too large.
Previous results for the microturbulence vary. GGGV82 found the same value in both stars, -14v t = 2.6 km s −1 , as did Gratton & Ortolani (1984), v t = 3.0 km s −1 . Gilroy et al. (1988) derived These comparisons are not very enlightening, but do suggest that the microturbulence values for the two stars are not very different.
This difference is not easily anticipated by past spectroscopic studies. Previously estimated for HD 115444 and -2.6 for HD 122563 (Bond 1980, uvby); -2.56 and -2.47 (Anthony-Twarog & Twarog 1994, uvby); -2.83 and -2.58 (Clariá et al. 1994;DDO δ4548). Thus taken as a whole, the previous spectroscopic and photometric investigations do not disagree with the modest metallicity difference found in the present work.
We performed a further check on the spectroscopically-derived gravity, by computing an "evolutionary" gravity using the standard formula log g evol = −12.50 + 0.4 (M V + BC) + log M + △ log T eff which combines Newton's gravitation law with the Stefan-Boltzmann law, eliminating the stellar radius R. We assumed masses equal to the standard mass of stars near the Population II turnoff, respectively. This is in good agreement with our log g estimates.
Hipparcos parallaxes (ESA 1997) exist for both stars: 3.55 ± 1.12 milliarcsec for HD 115444 and 3.76 ± 0.72 milliarcsec for HD 122563-almost the same values. But the apparent V of HD 122563 is 2.78 magnitudes smaller than that of HD 115444. The standard distance modulus relation (M V = V + 5 log p + 5 − A V ) then implies nearly the same difference in absolute magnitudes in the absence of significant ISM extinction. This is surprising since these stars, both red giants with very similar derived atmosphere parameters, would be expected to differ little in luminosity (as they indeed do according to literature estimates). Using the evolutionary gravity equation for HD 122563, and adopting Bond's (1980) estimate that A V = 0.0, we calculated from the Hipparcos parallax that log g Hipp = 1.42 ± 0.21, very close to our spectroscopic log g = 1.30.
The uncertainty in log g Hipp was computed from the maximum and minimum values permitted by the uncertainties in the quantities upon which it depends. Repeating the same calculation for HD 115444 (for which Bond also estimated A V = 0.0) gave log g Hipp = 2.58 ± 0.33; this is far from our estimate (in agreement with others) of log g = 1.50. From this discussion and the results of literature comparisons, probable statistical biases make it is very risky to use such small parallaxes with large relative errors (34% for HD 115444).

Light and Fe-peak Elements
With the final model atmospheres, we determined abundances for nearly all elements with Z ≤ 30 from their measured EWs. The exception was manganese, for which we treated the strong Mn i lines in synthetic spectrum calculations, obtaining the hyperfine substructure components from Kurucz's (1999) web site. Mean log ε(X) abundances and [X/Fe] relative abundance ratios are given for both stars in Table 4. For a species with at least three available transitions in a stellar spectrum, the standard deviation of the sample σ quoted in Table 4 is simply the internal scatter of individual measurements. The mean of those values is < σ > = 0.10 for HD 115444 (24 measures), and 0.12 for HD 122563 (15 measures). We take the average of these numbers as -16representative of the scatter of individual values of a given abundance for our stars. Thus for a species with only one transition we entered 0.11 as σ in Table 4, and for a species with two transitions we entered the greater of 0.11 and the measured σ. In subsequent discussions we will use these σ values for the abundance uncertainties, rather than the more commonly employed standard deviations of the means (σ/ √ #lines). Adoption of the more conservative σ allows many remaining uncertainties of scale (from choices of model atmosphere grids, transition probabilities, etc.) to be effectively accounted for. These σ's are also far more realistic abundance uncertainty estimates for species with small numbers of measured transitions.
In Figure 4 we However, some noteworthy anomalies exist in the abundance patterns. To see this more clearly, in Table 4 we also give for all species the abundance differences between the two stars. These differences, like those discussed in §3.1, were computed on a line-by-line basis. As before we define ∆ log ε(X) ≡ log ε(X) HD 115444 -log ε(X) HD 122563 . The single-measure standard deviation in the relative abundance σ(∆ log ε(X)) is usually much smaller than those of the absolute abundances σ(log ε(X)). For species with three or more transitions, < σ(∆ log ε(X))> = 0.06 ± 0.01 from 13 values, excluding the anomalously large scatter in Ba ii. Therefore 0.06 is entered as σ(∆ log ε(X)) for a species with only one transition, and the greater of 0.06 and the measured value is entered for a species with two transitions.
-17 -The abundance differences from Table 4 are displayed in Figure 5 (not every element was detected in HD 122563). For the elements with detected transitions of both neutral and ionized species (Ti, V, and Fe), we plot a single point representing the mean differences. The vast majority of elements with Z ≤ 30 (13 out of 18) have ∆ log ε(X) values within 2σ of ∆ log ε(Fe) = -0.24.
The abundances of these elements thus simply scale with the metallicity difference between the two stars. The most notable exceptions are the large relative C, Ti excesses and the significant relative Mn underabundance in HD 115444. For elements with Z > 30 our results, supporting and extending those of GGGV82 and Gilroy et al. (1982), show only a small enhancement of the lighter neutron-capture elements in HD 115444 but about a factor of ten difference among the heavier elements. We will now comment on these and other individual abundances: The CNO-group: carbon. The CNO elements were not of primary interest in this study, so to determine carbon and nitrogen abundances we simply adopted the atomic and molecular linelists from the Kurucz (1999) web site for synthetic spectrum analysis of selected portions of the CH A 2 ∆ − X 2 Π G-band and the CN B 2 Σ + − X 2 Σ + violet bands. With these linelists we derived a clear abundance enhancement of C in HD 115444 relative HD 122563 ( Fe-peak Elements. The Fe-peak abundance pattern in our stars (Figure 4) follows the general distribution in other UMP stars, and HD 115444 generally cannot be distinguished -19from HD 122563. A conspicuous exception to this statement is manganese, for which we derive very different abundances in the two stars: ∆ log ε(Mn) = -0.49, more than 4σ smaller than ∆ log ε(Fe). The Mn anomaly is easily seen from inspection of the data: whereas typical weak Fe-peak absorption features in the HD 115444 spectrum are about half as strong as those in the HD 122563 spectrum, the weak Mn i lines are a third as strong (  (Sneden, Gratton, & Crocker (1991).
Titanium (probably belonging to both α-capture and Fe-peak element groups) is overabundant, but much more so in HD 115444; the mean of Ti i and Ti ii relative abundances is ∆ log ε(Ti) = -0.06, which is about 3σ higher than ∆ log ε(Fe). As with the case of Mn, the Ti anomaly can be readily forecast by inspection of the spectra of the two stars. The Ti abundance difference is not easily explained, as it is based on large numbers of both neutral and ionized species transitions.
Departures from LTE, which are conceivable as such, probably cannot account for this abundance difference, since we must demand that any proposed non-LTE corrections should affect neutral and ionized species abundances in significantly different ways in the two stars.

The Neutron-Capture Elements
The abundances of the neutron-capture elements (Z > 30) were determined by EWs or by synthetic spectra, as indicated in Table 2. The linelists for these calculations were those originally developed for the abundance analysis of CS 22892-052 . Hyperfine structure was explicitly accounted for in the abundance computations for Ba, La, and Eu; see Sneden et al.
for further discussion and references to the laboratory data. An example of a synthetic spectrum fit to the same features in both stars is shown in Figure 6. It is clear that the derived abundance -20from the Eu ii feature in HD 115444 is greatly enhanced relative that in HD 122563. As HD 115444 is more metal-poor than HD 122563, the contrast between the europium abundance is very large between the two stars. This statement holds for all of the heavier (Z ≥ 56) elements (Table 4, Figure 5). On the other hand, the lighter neutron-capture elements (32 ≤ Z ≤ 40) are only weakly enhanced in HD 115444. These results were anticipated from the spectra presented in Figure 1, and are in qualitative agreement with the earlier studies of GGGV82 and Gilroy et al. (1988).
With this work we have increased the total number of neutron-capture element abundances in HD 115444 to 20. The reader is reminded that the abundances for Ge, Os, and Pt were obtained at λ2833.05Å. Our new inspection of that line re-affirmed the earlier concerns, and so we decided not to include Pb among the neutron-capture elements in our present work.

Thorium and Uranium
The Th ii λ4019.1Å line is the only thorium feature that has been employed for abundance determination in metal-poor stars, although recently Sneden et al. (1998a) have detected the weaker λ4086.5Å line in the extremely r-process rich star CS 22892-052. The λ4019Å line is heavily blended, and several previous studies (e.g., Morell et al. 1992, François et al. 1993, Norris et al. 1997b) have carefully considered possible contaminant lines to the total feature. A preliminary Th abundance in HD 115444 has been presented by Cowan et al. (1999a) from a spectrum combined from our 2d-coudé data and similar data from the Keck I HIRES. Here we discuss the analysis of just the 2d-coudé data.
-21 - We carried out synthetic spectrum computations for both stars, varying several elemental and CH abundances in order to get the best overall match to the observed feature. In Figure 7 we show the results of these experiments. In each panel, syntheses with different thorium abundances are displayed. The solid line represents the case where Th is deleted from the analyses. To set the strengths of the 13 CH lines, we first determined carbon isotope ratios from synthetic spectra of the A 2 ∆ − X 2 Π band lines near λ4230Å and λ4370Å, obtaining 12 C/ 13 C = 4.0 ±1.5 for HD 122563 and 6.0 ±1.5 for HD 115444 (in agreement with earlier studies by Sneden et al. 1986). Then we empirically determined the CH abundance that would best match the observed 12 CH B 2 Σ − X 2 Π lines surrounding λ4019Å (for example, the λλ4020.0,4020.2Å CH pair shown in Figure 7). Finally, the strengths of the 13 CH λλ4018.97,4019.14Å lines were scaled down by the 12 C/ 13 C ratio.

Discussion
Our abundance analyses have indicated (1) that the heavier neutron-capture elements are overabundant by a factor of 10 in HD 115444 compared to HD 122563.
(2) The iron-peak elements have similar abundances in HD 122563 and HD 115444, with the notable exceptions of Ti and Mn.
(3) The carbon abundance is high in HD 115444. We discuss these points in more detail in this section.

Comments on the CNO Abundances
Nearly all high luminosity (M V ∼ < 0) halo giant stars exhibit very low 12 C/ 13 C ratios and low total carbon abundances, indicating efficient operation of CN-cycle nucleosynthesis and convective envelope dredge-up along the giant branch. The required envelope mixing is much more than that predicted for the canonical first dredge-up in such low mass, low metallicity stars. The extra mixing almost surely begins at the red giant stage in which the outwardly advancing H-burning shell contacts the molecular weight barrier left behind by the first-dredge up. Recent theoretical studies of this phenomenon have been conducted by Charbonnel (1995); new observational confirmation and references to prior work are given by Gratton et al. (1999).
Our stars match well the general trends of CNO abundances in low metallicity giants. But the carbon abundance in HD 115444 relative to HD 122563 is large (∆ log ε(C) = +0.10, compared to ∆ log ε(Fe) = -0.24). Moreover, HD 122563 has a smaller isotopic ratio ( 12 C/ 13 C ≃ 4) than does HD 115444 ( 12 C/ 13 C ≃ 6), and it also has a smaller log ε(C). Therefore one obvious possibility is that HD 122563, at a more advanced giant branch evolutionary state (lower T eff , lower M V ) than HD 115444, has had a chance to mix more CN-cycle products to its surface. The At present, differences in convective envelope dredge-up efficiency seems the more likely explanation for the difference in carbon abundance of our two stars. We would like to ascribe the slightly higher 12 C/ 13 C ratio of HD 115444 to this phenomenon also, but the observational stars; see e.g. Laird 1985, Carbon et al. 1987, Tomkin & Lambert 1984, Tomkin, Sneden, & Lambert 1986 (Israelian et al. 1998, Boesgaard et al. 1999

Abundances of the Iron-Peak Elements
Recall from Figure 5 that the relative abundances ∆ log ε(X) of nearly all Fe-peak elements are mostly consistent with the relative iron abundance metallicity difference ∆ log ε(Fe) between the two stars. Given this unity (based on eight elements), the relative manganese abundance abundance ∆ log ε(Mn) ≃ -0.5 is quite different. Both stars are, however, within the range in The mild relative overabundance of Ti in HD 115444, ∆ log ε(Ti) = -0.06, also does not seem to be predicted from inspection of the figures in Nakamura et al.. The origin of this abundance difference may be due to differences in total (and/or core) masses of the probably relative small sample of supernovae that once produced the heavy elements in the two stars. This suggestion is supported by recent early galactic nucleosynthesis simulations by Karlsson & Gustafsson (1999).
Their models yield increasing scatter in [Ti/X], as well as in many other abundance ratios, as overall Fe metallicity decreases from [Fe/H] = -2.5 to -3.

Abundances of the Neutron-Capture Elements
The abundances of the elements Ba↔Dy in HD 115444 are significantly larger (∼ 10x) than in HD 122563, even though the iron metallicity is very similar in these stars (Table 4, Figure 5).
This gross difference between the absolute levels of the heaviest neutron-capture elements in these two stars is one more clear example that the abundances of the heavier elements vary greatly with respect to Fe in UMP stars of the early galaxy (Gilroy et al. 1988;Burris et al. 1999).
This presumably reflects the fact that the early Galaxy was not well mixed, with some stars receiving larger absolute amounts of neutron-capture material from individual nucleosynthetic events (almost certainly Type II SNe). Figure 8 shows the newly-determined neutron-capture element abundances for HD 115444 along with the HST abundances of Sneden et al. (1998a). Superimposed upon the abundance data is the scaled solar system r-process elemental abundance distribution. The solar system r-process curve has been scaled vertically downward to account for the lower metallicity of this star with respect to the Sun. This elemental curve was obtained by summing over the isotopic abundances, and these were based on deconvolving the total solar system abundances (Anders & Grevesse 1989) into s-and r-process abundance fractions, using the neutron capture cross sections of Käppeler et al. (1989) and Wisshak, Voss, & Käppeler (1996); see Sneden et al. (1996) and -26 - Burris et al. (1999) for details of this procedure.
The agreement between the heavier elements (Z ≥ 56) and the scaled solar system curve is very striking in Figure 8. While there has been increasing evidence that this pattern is true for elements in the range 56 ≤ Z ≤ 70, there have been few data for even heavier stable elements (76 ≤ Z ≤ 83, Os↔Bi, also know as 3 rd r-process peak). The new observations of HD 115444 confirm the scaled solar system r-process pattern throughout the range 56 ≤ Z ≤ 78.
This in turn argues for a common site for r-process synthesis of these elements. Possible sites include supernovae and neutron-star binaries; see Wheeler, Cowan, & Hillebrandt (1998) and Freiburghaus et al. (1999) for an extended theoretical discussion. But the observations will not easily support a wide range of allowable r-process parameters: a narrow range of progenitor star masses (e.g., Mathews, Bazan, & Cowan 1992;Wheeler et al. 1998) or a narrow range of nuclear and astrophysical conditions in the r-process sites (e.g., Freiburghaus et al. 1999) are required.
The lighter neutron-capture elements Ge, Sr, Y and Zr shown in Figure 5 present a puzzle, as their abundances are very similar in both stars. This suggests that the abundances of these elements are dependent upon metallicity, unlike the heavier elements (see also McWilliam et al. 1995, Sneden et al. 1998a, Burris et al. 1999, Cowan et al. 1999a. Moreover, these elements in both stars are poorly matched by the scaled solar system r-process curve in Figure 8. Further observational evidence supporting this point comes from the identification of several elements in the 41 ≤ Z ≤ 48 range from new near-UV spectra of CS 22892-052 (Cowan et al. 1999b). The scaled solar-system r-process curve is also not a good fit to the abundances of that more extensive set of elements with Z < 56. These mis-matches might demonstrate the existence of such a "second r-process" that synthesizes these lighter r-process nuclei in a different environment than the heavier r-process nuclei (see Wasserburg, Busso, & Gallino 1996;Qian, Vogel, & Wasserburg 1998). More extensive theoretical abundance predictions of these elements (see e.g., Gallino et al. 1999) are clearly needed here.
We have made quantitative comparisons of the abundances of elements with 56 ≤ Z ≤ 78 in HD 115444 and CS 22892-052 to the solar system r-process abundance distribution. Here, these -27elements will be collectively referred to as "heavy-r" elements. The weighted mean abundance offset of the heavy-r elements for HD 115444 is <[heavy-r/H]> = -2.06 ± 0.03 (σ = 0.10; 14 elements), and for CS 22892-052 <[heavy-r/H]> = -1.44 ± 0.02 (σ = 0.06; 16 elements). In these calculations, the weight employed for each element's abundance was its observed line-to-line scatter standard deviation. However, weighting by the number of lines used, or assigning equal weights to each element produced little change in the mean abundances. These means represent the r-process abundance levels of the heavy stable elements in the two stars. In the top and middle panels of Figure 9 we show the scatter about the means for HD 115444 and CS 22892-052.
The larger scatter for HD 115444 seen in the top panel is evidence of the weaker absorption lines of neutron-capture species in this star compared to those of CS 22892-052, which increased the analysis difficulties for some elements. For example, the Tm abundance for HD 115444 comes a single very weak and blended Tm ii transition at λ3700.2Å, while in CS 22892-052 four Tm ii lines were analyzed by Sneden et al. (1996). But in both stars it is clear that there is a single-valued abundance offset from the heavy-r solar system distribution.
The bottom panel of Figure 9 shows the comparison of a combined abundance set (HD 115444 plus CS 22892-052) with the heavy-r solar system values. The combined stellar set was formed by first normalizing each star's heavy-r abundances by the mean abundance described above (log ε(X) HD 115444 + 2.06, and log ε(X) CS 22892−052 + 1.44). Then we computed straight means of the normalized abundances for those elements analyzed in both stars, or else adopted the normalized abundances of one star when there was none available for the other star. Thus the abundance of terbium and hafnium in the combined set comes only from CS 22892-052, while the platinum abundance comes only from HD 115444. Comparing this normalized, combined data set to the solar-system r-process distribution yielded <[heavy-r/H]> = 0.00 ± 0.02 (σ = 0.08; 17 elements). No element abundance in the combined data deviates by more than 0.13 dex from the solar system distribution. Figure 9 clearly demonstrates that the abundances of the heavy stable elements in HD 115444 and CS 22892-052 from Z = 56-78 are consistent with the solar system r-process abundances. While this has been suggested previously for CS 22892-052 , we now have extensive neutron-capture element abundance data, including that in the 3 rd r-process peak, for two stars. These new data, particularly the normalized and combined abundances of HD 115444 and CS 22892-052 (as shown in Figure 9) yield solid agreement with solar values, making other explanations for this abundance pattern (e.g., Goriely & Arnould 1997;Goriely & Clerbaux 1999) less likely. (See also further discussion in Freiburghaus et al. 1999.) As more r-process-rich UMP stars are discovered and subjected to detailed analysis, the accidental errors should decrease even further, leaving any remaining deviations to mark either true departures from the solar abundances or residual abundance analysis problems.

Chronometers and Ages
The detection of radioactive thorium in HD 115444 (Figure 7) provides an opportunity to determine the radioactive age of this star. Consider first the chronometric age based solely on the HD 115444 abundances. To minimize errors in abundance determinations we have employed the ratio of thorium, produced solely in the r-process, to the stable abundance level of europium, an element overwhelmingly produced by the r-process (see Sneden et al. 1996, Burris et al. 1999).
While Cowan et al. (1999a) found a preliminary value of log ε(Th) = -2.1 ± 0.07 (based on an older model stellar atmosphere), our new synthesis indicates a value of log ε(Th) = -2.23 ± 0.07, for thorium and an observed ratio of log ε(Th/Eu) = -0.60 in HD 115444. This value is very similar to that found previously for CS 22892-052, of log ε(Th/Eu) = -0.66 .
The observed ratio of Th/Eu in HD 115444 can be compared to the initial predicted value at time of formation in an r-process environment. While the site for the r-process is still not well-defined, abundances of the heaviest stable neutron-capture elements detected in HD 115444 and other stars can be used to constrain the predictions of theoretical r-process calculations.
These same calculations can then be employed to predict the zero-decay age abundances of the radioactive elements Th and U. Recently Pfeiffer et al. (1997) and Cowan et al. (1999a), used the most massive stable solar isotopic abundances (i.e, 206,207,208 Pb and 209 Bi) to constrain the theoretical calculations. They conclude that the best fits to the observed solar data are obtained using the nuclear mass models FRDM-HFB (finite range droplet with Hartree-Fock-Bogoliubov), With the detailed observations for two stars at this point, our analysis suggests a reduced (with respect to only one star) observational error in log ǫ of ± 0.05. Applying that abundance error to the mean age determination suggests age uncertainties of +2.3 −2.4 Gyr, based solely upon observational errors. Extensive error analysis and discussion (Cowan et al. 1999a) suggests that the calculations best fitting the solar data have typical theoretical uncertainties of ± 3 Gyr. If we assume an average observational uncertainty of 2.4 Gyr, then the total error (added in quadrature -30for uncorrelated errors) is approximately ± 4 Gyr, resulting in a final value of 15.6 ± 4 Gyr for the age of these two UMP stars.
These calculations demonstrate both the power and the sensitivity of this method, with small changes in observed values or theoretical predictions changing the ages by the order of several Gyrs. While we have reduced the observational errors somewhat as a result of having two stars to analyze, it will be difficult to reduce further the total error without additional stellar observations of thorium and further theoretical calculations employing the most reliable mass formulae for nuclei far from stability. It is worth noting, however, that this radioactive-decay technique does not suffer from uncertainties in chemical evolution, since the time-scale for such evolution in the early Galaxy is much shorter than the radioactive decay half-life (14 Gyr) of Th and the age of the Galaxy.
Additional supporting age information could be provided by the uranium chronometer in HD 115444, but we were only able to set an upper limit on the abundance of this element of log ε(U) ≤ -2.6. Here is an age estimation based upon that upper limit. In contrast to the case for Th, in the case of U we have to take into account the two different isotopes, with very different half-lives, that make up this element. Theoretical calculations based upon the ETFSI-Q model (see Cowan et al. 1999a;Pfeiffer et al. 1997) predict that the initial U elemental abundance is dominated by the shorter half-life isotope, 235 U, while presumably the elemental abundance in the metal-poor stars is mostly determined by 238 U. Ignoring the initial 235 U production entirely (this isotope decays in much less than a Hubble time), and only comparing the initial time-zero 238 U and the observed upper limit in HD 115444 gives a lower limit on the age of this star of > 11 Gyr. This value is already within the estimated error for the Th age estimate. Since the actual U abundance may be much below our upper limit, the implied lower limit on the stellar age (based upon the non-detection of U) is in reasonable accordance with the Th age determination.
The thorium age estimates determined here, both for HD 115444 by itself and by considering both UMP stars, are consistent with our earlier age estimates for CS 22892-052, the (so far) only other metal-poor Galactic halo star for which abundances of Th and an extensive set of other -31neutron-capture elements have been determined (Cowan et al. 1997, Pfeiffer et al. 1997, Cowan et al. 1999a. Further, these new chronometric age determinations agree well with recent globular cluster age estimates (see Pont et al. 1998) and cosmological age determinations of 14.9 ± 1.5 Gyr (Perlmutter et al. 1999) and 14.2 ± 1.7 Gyr (Riess et al. 1998). While values of the Hubble constant near 70 km sec −1 Mpc −1 and a flat universe would imply cosmological age estimates much less than this, most recent work has suggested a relatively low density universe, which along with that Hubble constant have suggested a corresponding universal age of 13 Gyr (see e.g., Freedman 1999), still consistent with our estimated stellar chronometric ages. We note, however, that we have found thorium in only two stars so far, and additional detections will be needed to strengthen this technique of radioactive dating.

Conclusions
We have performed a detailed abundance analysis of the UMP field giant HD 115444. In addition to treating HD 115444 alone in a standard model atmosphere analysis of equivalent widths and synthetic spectra, we have made a line-by-line abundance comparison with the well-known UMP giant HD 122563. The differential study yields very accurate relative atmosphere parameters and abundances that illuminate a number of similarities and sharp disagreements in the chemical compositions of these two stars. Here we summarize the main results, which confirm and greatly extend the initial spectroscopic investigations by GGGV82 and Gilroy et al. (1988).
• HD 115444's mean Fe-peak metallicity is approximately 0.24 dex lower than that of HD 122563. It is a slightly less evolved star, with ∆ T eff ≃ +150 K, ∆ log g ≃ +0.2, and ∆ M V ≃ +0.6. The absolute metallicity of HD 115444 is [Fe/H] = -2.99, in good agreement with GGGV82, but this number is somewhat dependent on the adopted temperature scale, choice of model atmosphere, and assumptions made in constructing the models.
• For elements with Z ≤ 30, HD 115444 has a fairly typical halo star abundance distribution.
But comparison to HD 122563 reveals several significant differences in abundance patterns.
-32 -Most notable are the large relative overabundance of carbon and underabundance of manganese in HD 115444, and the smaller but significant overabundance of titanium.
• Thorium is underabundant relative to the other heavy neutron-capture elements, as it is in CS 22892-052. A simple age estimate for these two stars is 15.6 ± 4 Gyr.
We thank Roberto Gallino, Inese Ivans, Francesca Primas, Friedel Thielemann, Jim Truran, Craig Wheeler and referee Roger Cayrel for helpful discussions and comments on this paper.