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Register a new userImproved Atomic Transition Probabilities for UV and Optical Lines of Hf II and Determination of the Hf Abundance in Two Metal-Poor Stars
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We report new branching fraction measurements for 199 UV and optical transitions of Hf II. These transitions range in wavelength (wavenumber) from 2068- 6584 A (48322-15183 cm-1) and originate in 17 odd-parity upper levels ranging in energy from 38578-53227 cm-1. The branching fractions are combined with radiative lifetimes reported in an earlier study to produce a set of transition probabilities and log(gf) values with accuracy ranging from 5-25%. Comparison is made to transition probabilities from the literature where such data exist. We use these new transition probabilities to derive improved Hf abundances in two metal-poor stars. HD 196944 is enhanced in s-process elements, and we derive log epsilon (Hf) = -0.72 +/- 0.03 (sigma = 0.09) from 12 Hf II lines. HD 222925 is enhanced in r-process elements, and we derive log epsilon (Hf) = 0.32 +/- 0.03 (sigma = 0.11) from 20 Hf II lines. These measurements greatly expand the number of potentially useful Hf II lines for analysis in UV and optical spectra.
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We report new branching fractions for 121 UV lines from the low-lying odd-parity levels of Fe II belonging to the z6Do, z6Fo, z6Po, z4Fo, z4Do and z4Po terms of the 3d6(5D)4p configuration. These lines range in wavelength from 2250 - 3280 {AA} and originate in levels ranging in energy from 38459 - 47626 cm-1. In addition, we report branching fractions for 10 weak blue lines connecting to the z4Do term which range in wavelength from 4173 - 4584 {AA}. The BFs are combined with radiative lifetimes from the literature to determine transition probabilities and log(gf) values. Comparison is made to selected experimental and theoretical data from the literature. Our new data are applied to iron abundance determinations in the Sun and in metal-poor star HD 84937. For the Sun, eight blue lines yield log {epsilon}(Fe) = 7.46 +/- 0.03, in agreement with standard solar abundance estimates. For HD 84937 the observable wavelength range extends to the vacuum UV ({lambda} >= 2327 {AA}), and from 75 lines we derive log {epsilon}(Fe) = 5.26 +/- 0.01 ({sigma} = 0.07), near to the metallicity estimates of past HD 84937 studies.
(Abridged) Extremely metal-poor stars contain the fossil records of the chemical composition of the early Galaxy. The NLTE profiles of the calcium lines were computed in a sample of 53 extremely metal-poor stars with a modified version of the program MULTI. With our new model atom we are able to reconcile the abundance of Ca deduced from the Ca I and Ca II lines in Procyon. -We find that [Ca/Fe] = 0.50 $pm$ 0.09 in the early Galaxy, a value slightly higher than the previous LTE estimations. -The scatter of the ratios [X/Ca] is generally smaller than the scatter of the ratio [X/Mg] where X is a light metal (O, Na, Mg, Al, S, and K) with the exception of Al. These scatters cannot be explained by error of measurements, except for oxygen. Surprisingly, the scatter of [X/Fe] is always equal to, or even smaller than, the scatter around the mean value of [X/Ca]. -We note that at low metallicity, the wavelength of the Ca I resonance line is shifted relative to the (weaker) subordinate lines, a signature of the effect of convection. -The Ca abundance deduced from the Ca I resonance line (422.7 nm) is found to be systematically smaller at very low metallicity, than the abundance deduced from the subordinate lines.
Recent radiative lifetime measurements accurate to +/- 5% using laser-induced fluorescence (LIF) on 43 even-parity and 15 odd-parity levels of Ce II have been combined with new branching fractions measured using a Fourier transform spectrometer (FTS) to determine transition probabilities for 921 lines of Ce II. This improved laboratory data set has been used to determine a new solar photospheric Ce abundance, log epsilon = 1.61 +/- 0.01 (sigma = 0.06 from 45 lines), a value in excellent agreement with the recommended meteoritic abundance, log epsilon = 1.61 +/- 0.02. Revised Ce abundances have also been derived for the r-process-rich metal-poor giant stars BD+17 3248, CS 22892-052, CS 31082-001, HD 115444 and HD 221170. Between 26 and 40 lines were used for determining the Ce abundance in these five stars, yielding a small statistical uncertainty of 0.01 dex similar to the Solar result. The relative abundances in the metal-poor stars of Ce and Eu, a nearly pure r-process element in the Sun, matches r-process only model predictions for Solar System material. This consistent match with small scatter over a wide range of stellar metallicities lends support to these predictions of elemental fractions. A companion paper includes an interpretation of these new precision abundance results for Ce as well as new abundance results and interpretations for Pr, Dy and Tm.
Our current understanding of the chemical evolution of the Universe is that a first generation of stars was formed out of primordial material, completely devoid of metals (Pop III stars). This first population of stars comprised massive stars that exploded as supernovae disseminating the metals they synthesised in the interstellar medium. These massive stars are long dead and cannot be observed in the local Universe. Among very metal poor stars (metallicity below -2.0) we expect to find the direct descendants of these pristine metal factories. The chemical composition of these stars provides us indirect information on the nature of the Pop III stars, their masses, luminosities and mode of explosion. The constraints are stronger if the chemical inventory is more complete, more chemical elements and isotopic ratios are measured for each star. Unfortunately the lower the metallicity of the star, the weaker the lines. Access to the space UV spectral range gives us crucial supplementary information. To start with, it allows access to some very strong Fe lines that may allow to measure the abundance of this element in stars for which this was not possible from the ground-accessible UV spectra. The number of such stars is steadily increasing. Next the UV range allows us to measure elements that cannot be measured from ground-based spectra like P, Ge, As, Se, Cd, Te, Lu, Os, Ir, Pt, Au. In addition it is fundamental for measuring other elements that can be accessed from earth, but with great difficulty, like C, S, Cu, Zn, Pb. The Hubble space telescope, with its limited collecting power made this possible only for very few stars. Old metal poor stars are cool, of spectral types F,G,K, and their UV flux is low. The availability of a UV high resolution spectrograph fed by a large area space telescope will open an unprecedented window on the early evolution of our Galaxy.
We present vanadium (V) abundances for 255 metal-poor stars, derived from high-resolution optical spectra from the Magellan Inamori Kyocera Echelle spectrograph on the Magellan Telescopes at Las Campanas Observatory, the Robert G. Tull Coud{e} Spectrograph on the Harlan J. Smith Telescope at McDonald Observatory, and the High Resolution Spectrograph on the Hobby-Eberly Telescope at McDonald Observatory. We use updated V I and V II atomic transition data from recent laboratory studies, and we increase the number of lines examined (from 1 to 4 lines of V I, and from 2 to 7 lines of V II). As a result, we reduce the V abundance uncertainties for most stars by more than 20% and expand the number of stars with V detections from 204 to 255. In the metallicity range $-$4.0 $<$ [Fe/H] $< -$1.0, we calculate the mean ratios [V I/Fe I]$ = -0.10 pm 0.01 (sigma = 0.16)$ from 128 stars with $geq$ 2 V I lines detected, [V II/Fe II] $= +0.13 pm 0.01 (sigma = 0.16)$ from 220 stars with $geq$ 2 V II lines detected, and [V II/V I] $= +0.25 pm 0.01 (sigma = 0.15)$ from 119 stars. We suspect this offset is due to non-LTE effects, and we recommend using [V II/Fe II], which is enhanced relative to the solar ratio, as a better representation of [V/Fe]. We provide more extensive evidence for abundance correlations detected previously among scandium, titanium, and vanadium, and we identify no systematic effects in the analysis that can explain these correlations.
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