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Astronomy & Astrophysics manuscript no. reference2 June 29, 2009

c ESO 2009

The chemical abundance analysis of normal early A- and late B-type stars
L. Fossati1 , T. Ryabchikova1,2 , S. Bagnulo3 , E. Alecian4,5 , J. Grunhut5 , O. Kochukhov6 , and G. Wade
1 2 3 4 5 6

5

¨ ¨ ¨ Institut fur Astronomie, Universitat Wien, Turkenschanstrasse 17, 1180 Wien, Austria. e-mail: fossati@astro.univie.ac.at,ryabchik@astro.univie.ac.at Institute of Astronomy, Russian Academy of Sciences, Pyatnitskaya 48, 119017 Moscow, Russia. e-mail: ryabchik@inasan.ru Armagh Observatory, College Hill, Armagh BT61 9DG, Northern Ireland, UK. e-mail: sba@arm.ac.uk Observatoire de Paris-Meudon, LESIA, UMR 8111 du CNRS, 92195 Meudon Cedex, France. e-mail: evelyne.alecian@obspm.fr Physics Dept., Royal Military College of Canada, PO Box 17000, Station Forces, K7K 4B4, Kingston, Canada. e-mail: Jason.Grunhut@rmc.ca,Gregg.Wade@rmc.ca Department of Physics and Astronomy, Uppsala University, SE-751 20, Uppsala, Sweden. e-mail: Oleg.Kochukhov@fysast.uu.se
ABSTRACT

Context. Modern spectroscopy of early-type stars often aims at studying complex physical phenomena such as stellar pulsation, the

peculiarity of the composition of the photosphere, chemical stratification, the presence of a magnetic field, and its interplay with the stellar atmosphere and the circumstellar environment. Comparatively less attention is paid to identifying and studying the "normal" A- and B-type stars and testing how the basic atomic parameters and standard spectral analysis allow one to fit the observations. By contrast, this kind of study is paramount eventually for allowing one to correctly quantify the impact of the various physical processes that occur inside the atmospheres of A- and B-type stars. Aims. We wish to establish whether the chemical composition of the solar photosphere can be regarded as a reference for early Aand late B-type stars. Methods. We have obtained optical high-resolution, high signal-to-noise ratio spectra of three slowly rotating early-type stars (HD 145788, 21 Peg and Cet) that show no obvious sign of chemical peculiarity, and performed a very accurate LTE abundance analysis of up to 38 ions of 26 elements (for 21 Peg), using a vast amount of spectral lines visible in the spectral region covered by our spectra. Results. We provide an exhaustive description of the abundance characteristics of the three analysed stars with a critical review of the line parameters used to derive the abundances. We compiled a table of atomic data for more than 1100 measured lines that may be used in the future as a reference. The abundances we obtained for He, C, Al, S, V, Cr, Mn, Fe, Ni, Sr, Y, and Zr are compatible with the solar ones derived with recent 3D radiative-hydrodynamical simulations of the solar photosphere. The abundances of the remaining studied elements show some degree of discrepancy compared to the solar photosphere. Those of N, Na, Mg, Si, Ca, Ti, and Nd may well be ascribed to non-LTE effects; for P, Cl, Sc and Co, non-LTE effects are totally unknown; O, Ne, Ar, and Ba show discrepancies that cannot be ascribed to non-LTE effects. The discrepancies obtained for O (in two stars) and Ne agree with very recent non-LTE abundance analysis of early B-type stars in the solar neighbourhood.

1. Introduction
In the last decade there has been dramatic improvement in the tools for the analysis of optical stellar spectra, both from the observational and theoretical perspective. New high-resolution echelle instruments have come online, which cover much broader spectral ranges than older single-order spectrographs. Data quality has also substantially improved in terms of signalto-noise ratio (SNR), because of substantially greater instrument efficiency, and the use of large-size telescopes. Thanks to the vibrant observational activities of the past few years, and thanks to efficient and user-friendly data archive facilities, a huge highquality spectroscopic database is now available to the public. With the development of powerful and cheap computers, it has become practical to exploit these new data by performing spectral analysis using large spectral windows rather than selected spectral lines, at a level of realism heretofore impos-

sible. The high accuracy of observations and modelling techniques now allows, for example, stretching the realm of abundance analysis to faster rotators than was possible in the past, but also provides the possibility of learning more about the structure of the stellar atmospheres and the ongoing physical processes, especially when spectral synthesis fails to reproduce the observations. For instance, observed discrepancies between observed and synthetic spectra have allowed us to discover that the signature of chemical stratification is ubiquitous in the spectra of some chemically peculiar stars (Bagnulo et al. 2001; Wade et al. 2001; Ryabchikova et al. 2003) and to perform accurate modelling of this stratification in the atmospheres of Ap stars (for instance, Ryabchikova et al. 2005; Kochukhov et al. 2006). However, our inability to reproduce observations frequently stems for a very simple cause: that atomic data for individual spectral lines are incorrect. For solar type stars it is possible to construct a "reference" list of reliable spectral lines with reliable


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Fossati et al.: The abundance analysis for normal A- and B-type stars.

atomic data through the comparison of synthetic spectra with the solar observed spectrum, since the solar abundances are accurately known. In many cases, however, the solar spectrum cannot provide the required information, because the temperature of the target stars is significantly different from that of the sun. This problem can be overcome by adopting an analogous reference at a temperature reasonably close to the target temperature. The method involves a selection of a set of suitable reference stars for which very high quality spectra are available. Then an accurate determination of the stellar photospheric parameters and an accurate abundance analysis are performed with the largest possible number of spectral lines and the best possible atomic data. Finally, those spectral lines exhibiting the largest discrepancies from the model fit are identified, and their atomic data revised by assuming that the average abundance (inferred from the complete sample of spectral lines of that element) is the correct one. In this process it is important to take effects into account that can potentially play a significant role in all stellar atmospheres, such as variations in the model structure from non-solar abundances and non-LTE effects. In this paper we address the problem of establishing references for effective temperatures around 10000­13000 K. This is in some respect the easiest temperature range to study, as well as one of the most interesting. This temperature is close to ideal because the spectra of stars in this interval are generally unaffected by severe blending. It is also relevant because stars in this temperature range display spectroscopic peculiarities (chemical abundance peculiarities, stratification, Zeeman effect, etc.) that reflect physical conditions and processes of interest for detailed investigation. A crucial prerequisite for studying and characterising these phenomena is the capacity to model the underlying stellar spectrum in detail, and this requires high quality atomic data. The highest degree of accuracy in abundance analysis is reached for sharp-lined stars. Unfortunately, these objects are quite rare among A- and B-type stars, which are generally characterised by high rotational velocities. Furthermore, most of the slowly rotating stars in the chosen temperature region belong to various groups of magnetic and non-magnetic, chemically peculiar objects. As a matter of fact, many previous studies aimed at determining the chemical composition of "normal" early A- and late B-type stars were based on samples "polluted" by moderately chemically peculiar stars. For instance, the work by Hempel & Holweger (2003) includes the sharp-lined HgMn star 53 Aur. Hill & Landstreet (1993) searched for compositional differences among A-type stars, but four out of six programme stars are in fact classified as hot Am stars on the basis of the abundances of the heavy elements Sr-Y-Zr-Ba, which are considered as diffusion indicators (i.e., Sirius, o Peg, and Leo, see Hempel & Holweger 2003). The complexity of the problem of distinguishing between normal A and marginal Am stars is further stressed by Adelman & Unsuree (2007). The aim of the present paper is to search for sharp-lined early A- or late Btype stars with a chemical composition as close as possible to the solar one. As a final outcome, one could assess whether the chemical composition of the solar photosphere may be considered at least in principle as a reference for the composition of the early A- and late B-type stars. If such a star is found, this will not imply that the solar composition is the most characteristic for the slowly rotating A- and B-type stars, but will be used as further evidence that any departure from the composition of the solar photosphere has to be explained in terms of diffusion or other physical mechanisms that are not active at the same efficiency level in the solar photosphere.

Our work is based on a very detailed and accurate study of a vast sample of spectral lines. As a by-product, we provide a list of more than 1100 spectral lines from which we have assessed the accuracy of the corresponding atomic data. Such a list may serve as a future reference for further abundance analysis studies of stars with a similar spectral type. This paper is organised as follows. Sect. 2 describes the target selection, observations and data reduction, Sect. 3 presents methods and results for the choice of the best fundamental parameters that describe the atmospheres of the programme stars, Sect. 4 presents the methods and results of the abundance analysis of the programme stars. Our results are finally discussed in Sect. 5.

2. Target selection, observations, and data reduction
For our analysis we need a late A-type or early B-type star with a sharp-line spectrum (hence the star must have a small sin i) and one exhibiting the least possible complications due to phenomena such as non-homogeneous surface distribution of chemical elements, pulsation, or a magnetic field. Finding such a target is not a simple task, because most of the early-type stars are fast rotators. Slowly rotating A- and B-type stars generally show some type of chemical peculiarity often associated to the presence of abundance patches, a magnetic field (which broadens, or even splits spectral lines), and chemical stratification. Even Vega, which has been considered for a long time as the prototype of a "normal", slow rotating A-type star, is in fact currently classified as a Boo star and discovered to actually be a fast rotating star seen pole-on, exhibiting, as such, distorted line profiles (Adelman & Gulliver 1990; Yoon et al. 2008). Based on our knowledge, we reached the conclusion that the most suitable target for our project is the B9 star 21 Peg (HD 209459), which is known from previous studies as a "normal" single star with sin i 4 km s-1 (see, e.g. Sadakane 1981). We felt it was necessary to consider additional targets of our spectral analysis, for two main reasons. Since we intend to provide an accurate reference for the typical abundances of the chemical elements in A- and B-type stars (and compare these values with the solar ones), we need to cross-check with further examples whether the results obtained for 21 Peg are similar to those of other "normal", slow rotating A-type stars. Second, to check the accuracy of the astrophysical measurements of the logg f values, which is a natural complement of the present work. Both these goals are best achieved with the use of abundance values that have been obtained with a homogeneous method, rather than from a mixed collection of data from the literature. Therefore, we have also analysed another two stars of similar temperature as 21 Peg, i.e., HD 145788 (HR 6041) and Cet (HR 811). Both stars fulfill our requirements, although are slightly less ideal than 21 Peg. HD 145788, suggested to us by Prof. Fekel, is a slowly rotating single star with sin i 8 km s-1 (Fekel 2003). Cet, a SB1 with sin i 20 km s-1 , shows an infrared excess, and is a suspected Herbig Ae/Be star (Malfait et al. 1998). Since its spectrum is not visibly contaminated by the companion, it still serves our purpose. Cet was also already used as a normal comparison star in the abundance study of chemically peculiar stars by Smith & Dworetsky (1993). The star 21 Peg was observed five times during two observing nights in August 2007, with the FIES instrument of the North Optical Telescope (NOT). FIES is a cross-dispersed high-


Fossati et al.: The abundance analysis for normal A- and B-type stars.

3

´ resolution echelle spectrograph that offers a maximum spectral resolution of R = 65 000, covering the entire spectral range 3700­7400 å. Data were reduced using a pipeline developed by D. Lyashko, which is based on the one described by Tsymbal et al. (2003). All bias and flat-field images were median-averaged before calibration, and the scattered light was subtracted by using a 2D background approximation. For cleaning cosmic ray hits, an algorithm that compares the direct and reversed observed spectral profiles was adopted. To determine the boundaries of echelle orders, the code uses a special template for each order position in each row across the dispersion axis. The shift of the row spectra relative to the template was derived by a crosscorrelation technique. Wavelength calibration of each image was based on a single ThAr exposure, recorded immediately after the respective stellar time series, and calibrated by a 2D approximation of the dispersion surface. An internal accuracy of 100 ms-1 was achieved by using several hundred ThAr lines in every echelle order. Each reduced spectrum has a SNR per pixel of about 300 at 5000 å. All five spectra are fully consistent among themselves, which confirms that the star is not variable. This allowed us to combine all data in a unique spectrum with a final SNR of about 700. Because of the very low sin i of 21 Peg, we made use of a very high-resolution spectrum (R = 120 000) obtained with the Gecko instrument (now decommissioned) of the CanadaFrance-Hawaii Telescope (CFHT, Landstreet 1998) to measure this parameter. The spectrum covers the ranges 4612­4640 å and 5160­5192 å, which are too short to perform a full spectral analysis, but sufficient to measure sin i with high accuracy. According to Hubrig et al. (2006) the mean longitudinal magnetic field of 21 Peg is -144±60 G, which excludes the possibility that the star has a structured magnetic field. The spectrum of HD 145788 was obtained with the cross´ dispersed echelle spectrograph HARPS instrument attached at the 3.6-m ESO La Silla telescope, with an exposure time of 120 s, and reduced with the online pipeline 1 . The reduced spectrum has a resolution of 115 000, and a SNR per pixel of about 200 at 5000 å. The spectral range is 3780­6910 å with a gap between 5300 å and 5330 å, because one echelle order is lost in the gap between the two chips of the CCD mosaic detector. The star Cet was observed with the ESPaDOnS instrument of the CHFT February, 20 and 21 2005. ESPaDOnS consists of a table-top, cross-dispersed echelle spectrograph fed via a double optical fiber directly from a Cassegrain-mounted polarisation analysis module. Both Stokes I and V spectra were obtained throughout the spectral range 3700 to 10400 å at a resolution of about 65 000. The spectra were reduced using the Libre-ESpRIT reduction package Donati et al. (1997, and in prep.). The two spectra (each obtained from the combination of four 120 s subexposures) were combined into a final spectrum that has an SNR per pixel of about 1200 at 5000 å. The observation of Cet enters in the context of a large spectropolarimetric survey of Herbig Ae/Be stars. Least-squares deconvolution (LSD, Donati et al. 1997) was applied to the spectra of Cet assuming a solar abundance line mask corresponding to an effective temperature of 13000 K. The resulting LSD profiles show a clean, relatively sharp mean Stokes I profile, corresponding to sin i = 20 ± 1 km s-1 , and no detection of
1 http://www.ls.eso.org/lasilla/sciops/3p6/harps/ software.html#pipe

any Stokes V signature indicative of a photospheric magnetic field. Integration of Stokes V across the line using Eq. (1) of Wade et al. (2000) yields longitudinal magnetic fields consistent with zero field and with formal 1 uncertainties of about 10 G. The high-resolution spectropolarimetric measurements therefore provide no evidence of magnetic fields in the photospheric layers of the star. All the spectra of the three stars were normalised by fitting a spline to carefully selected continuum points. For each object, radial velocities r were determined by computing the median of the results obtained by fitting synthetic line profiles of several individual carefully selected lines into the observed spectrum. The r values are listed in Table 1, and their uncertainty is of the order of 0.5 km s-1 . The spectra were then shifted to the wavelength rest frame. Selected spectral windows containing the observed blue He I lines, together with the synthetic profiles, are displayed in Fig. 1.

HD 145788 1.4 1.2

normalised flux

1 0.8 0.6

21 Peg

Cet

0.4 0.2

4385

4390

4470

4475
o

4920

4924

wavelength (A)
Fig. 1. Samples of the three He I lines: 4387 final synthetic profiles tra of HD 145788 and respectively. spectra of HD 145788, 21 Peg and Cet around å, 4471 å and 4921 å in comparison with our (dashed lines) calculated for each line. The spec Cet are shifted upwards and downwards of 0.5,

3. Fundamental parameters
Fundamental parameters for the atmospheric models were obtained using photometric indicators as a first guess. For their refined estimate, we performed a spectroscopic analysis of hydrogen lines and metal lines, and as a final step, compared the observed and computed energy distributions. The spectroscopic tuning of the fundamental parameters is needed since different photometries and calibrations would give different parameters and uncertainties. The spectroscopic analysis will provide a set of parameters that fit all the parameter indicators consistently, with less uncertainties. Model atmospheres were calculated with LLMODELS, an LTE code that uses direct sampling of the line opacities (Shulyak et al. 2004) and allows computing models with an individualised abundance pattern. Atomic parameters of spectral lines used for model atmosphere calculations were extracted from the VALD database (Piskunov et al. 1995; Kupka et al. 1999; Ryabchikova et al. 1999).


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Fossati et al.: The abundance analysis for normal A- and B-type stars.

Before applying the spectroscopic method, we estimated the star's sin i. For 21 Peg, a sin i value of 3.76 ± 0.35 km s-1 was derived from the fit with a synthetic spectrum to 21 carefully selected lines observed with the Gecko instrument. This value agrees very well with the 3.9 km s-1 value derived by Landstreet (1998). Achieving such high precision was possible thanks to the high quality of the spectrum and the low sin i. The sin i values for HD 145788 and Cet, 10.0 ± 0.5 km s-1 and 20.2 ± 0.9 km s-1 , respectively, were derived from fitting about 20 wellselected lines along the whole available spectral region. In the next sections, we describe the determination of the atmospheric parameters: T eff , effective gravity (log g), and microturbulent velocity (mic ), and their uncertainties. The fundamental parameters finally adopted for 21 Peg, HD 145788, and Cet are given in Table 1.
3.1. Photometric indicators

Since none of the three programme stars, HD 145788, 21 Peg or Cet are known to be photometrically variable or peculiar, we can use temperature and gravity calibrations of different photometric indices for normal stars to get a preliminary estimate of the atmospheric parameters. The effective temper¨ ature (T eff ) and gravity (log g) were derived from Stromgren photometry (Hauck & Mermilliod 1998) with calibrations by Moon & Dworetsky (1985) and by Napiwotzki et al. (1993), and from Geneva photometry (Rufener 1988) with the calibration by North & Nicolet (1990). The mean parameters from the three calibrations that were used as starting models are the following ones: T eff = 9675±75 K, log g = 3.72±0.03 for HD 145788; T eff = 10255±115 K, log g = 3.51 for 21 Peg; T eff = 13200±65 K log g = 3.77±0.15 for Cet. No error bar is given for the log g of 21 Peg since all three calibrations give the same value.
3.2. Spectroscopic indicators 3.2.1. Hydrogen lines

For a fully consistent abundance analysis, the photometric parameters have to be checked and eventually tuned according to spectroscopic indicators, such as hydrogen line profiles. In the temperature range where HD 145788, 21 Peg, and Cet lie, the hydrogen line wings are less sensitive to T eff than to log g variations, but temperature effects can still be visible in the part of the wings close to the line core. For this reason hydrogen lines are very important not only for our analysis, but in general for every consistent parameter determination. To spectroscopically derive the fundamental parameters from hydrogen lines, we fitted synthetic line profiles, calculated with SYNTH3 (Kochukhov 2007), to the observed profiles. SYNTH3 incorporates the code by Barklem et al. (2000)2 that takes into account not only selfbroadening but also Stark broadening (see their Sect. 3). For the latter, the default mode of SYNTH3, adopted in this work, uses an improved and extended HLINOP routine (Kurucz 1993). Figure 2 shows the comparison between the observed H line profile for 21 Peg and the synthetic profiles calculated with the adopted stellar parameters. In Fig. 2 we also added the synthetic line profiles calculated with 1 error bars on T eff (± 200 K; upper profile) and log g (± 0.1 dex; lower profile). The same profiles with the same uncertainties, but for H, H, and H (from left to right) for the three programme stars, are shown in Figs. 9, 8, and 10 (online material).
2

http://www.astro.uu.se/barklem/hlinop.html

The three hydrogen lines (H, H, and H) used to spectroscopically improve the fundamental parameters for HD 145788 gave slightly different results both for T eff and for log g. As final values, we took their mean (T eff = 9750 K; log g = 3.7). This discrepancy is visible in Fig. 9 (online material), but it lies within the errors given for T eff and log g. The spectrum of HD 145788 also allowed a good normalisation of the region bluer than H. We synthesised this region to check the quality of our final fundamental parameters finding a very good fit for the three hydrogen lines H, H , and H8 . For 21 Peg we obtained the same temperature estimates from H and H (T eff = 10400 K) and by 200 K less from the fitting of H. We adopted a final T eff of 10400 K. To fit all three hydrogen lines, we need slightly different values of log g: 3.47 for H, 3.54 for H and 3.57 for H. We finally adopted log g = 3.55 taking possible continuum normalisation problems into account, in particular, for H. The temperature determination for Cet was more difficult thanks to the weak effect that this parameter has on the hydrogen lines at about 13000 K and to the slightly peculiar shape of the H line. As explained in Sect. 5, Cet probably shows a small emission signature in the region around the core of H possibly because of a circumstellar disk. This region is the one where T eff effects are visible, making it almost impossible to obtain a good temperature determination from this line. H and H gave best temperatures of 12700 K and 12900 K, respectively, leading to a final adopted value of 12800 K. Confirmation of this value was then given by the spectrophotometry (see Sect. 3.3). The results of LTE abundance analysis (Sect. 4) show that Cet has little He overabundance that leads to an overestimation of log g if the He abundance is not taken into account in the model atmosphere calculation (Auer et al. 1966). For this reason we derived the first set of fundamental parameters (T eff = 12800 K; log g = 3.80) and then the He abundance (log(NHe /Ntot ) = -0.97 dex). As a next step we recalculated a set of model atmospheres with the derived He abundance and re-fit the hydrogen line profiles. The best-fit gave us the same temperature, but weak effective gravity (log g = 3.75). As expected, the He overabundance is acting as pressure, requiring an adjustment of log g to be balanced. We obtained the He abundance from the fitting of the blue He line wings. The blue He lines are, in general, considered as showing very little non-LTE effect (Leone & Lanzafame 1998), and as we only used the line wings, this leads us to conclude that our results should not be affected by non-LTE effects and that the He is overabundant in Cet. The best fit to the blue He I lines of Cet is shown in Fig. 1. The example of Cet is important because it clearly shows the effect of the single element abundance on the parameter determination, not only for chemically peculiar stars (for which this effect is well known and often, but not always, taken into account), but also for chemically "normal" stars. The set of parameters that best fit the hydrogen line profiles could not be unique. For 21 Peg, we checked that using a combination of a lower temperature and lower gravity or else higher temperature and higher gravity increases the standard deviation of the fit of the H line wings by 25%. The result is that a different combination of T eff and log g could in principle provide a similar fit. For this reason the derived fundamental parameters should be checked with other indicators, such as the analysis of metallic lines (ionisation and excitation equilibrium) and the fitting of the spectral energy distribution. The latter is more important because ionisation and excitation equilibrium should be strictly used only with a full non-LTE treatment of the line formation.


Fossati et al.: The abundance analysis for normal A- and B-type stars. Table 1. Adopted atmospheric parameters for the analysed stars. Star Name 21 Peg HD 145788 Cet T eff [K] 10400 9750 12800 log g [cgs] 3.55 3.70 3.75 mic [km s-1 ] 0.5 1.3 1.0 mic [km s-1 ] 0.4 0.2 0.5 sin i [km s-1 ] 3.76 10.0 20.2 sin i [km s-1 ] 0.35 0.5 0.9 r [km s-1 ] 0.5 -13.9 12.5

5

The uncertainties on T eff , log g, and r are 200 K, 0.1 dex, and 0.5 km s-1 , respectively.

3.2.2. Metallic lines

The metallic-line spectrum may also provide constraints on the atmospheric parameters. If no deviation from the local thermodynamic equilibrium (LTE) is expected, the minimisation of the correlation between individual line abundances and excitation potential, for a certain element/ion, allows one to check the determined T eff . Then the balance between different ionisation stages of the same element provides a check for log g. The microturbulent velocity, mic is then determined by minimising the correlation between individual abundances and equivalent widths for a certain element. Determining the fundamental parameters performed in this way has to be done iteratively since, for example, a variation in T eff leads to a change in the best log g and mic .
-4

-4.2 -4.4 -4.6 -4.8 -5 10 20 30 40 50 60 70 80
o

Fe II Fe I 90 100 110

equivalent width (mA) Fe abundance
-4 -4.2 -4.4 -4.6 -4.8 -5 1 2 3 4 5 6 7 8 9 10 11

The uncertainties listed in Table 1 are given considering 2 on the error bar of the derived slopes. Considering a 1 error bar, the uncertainties on mic are of 0.1 km s-1 for HD 145788 and 21 Peg and of 0.2 km s-1 for Cet. According to previous works, deviation from LTE of the Fe II lines is expected to be very small ( 0.02 dex Gigas 1986; Rentzsch-Holm 1996) for the analysed stars, such that the absence of any clear correlation between the line Fe II abundance and the excitation potential confirms the T eff derived from the hydrogen lines. For the Fe I lines, deviations from LTE of about +0.3 dex are given by Gigas (1986) and Rentzsch-Holm (1996). However, both Gigas (1986) and Rentzsch-Holm (1996), as well as Hempel & Holweger (2003), used the same model atom, which includes 79 Fe I and 20 Fe II energy levels. We note that the highest energy level in their model atom for Fe II has an excitation energy of about 6 eV, while the ionisation potential is 16.17 eV. Such a model atom does not provide collisional coupling of Fe II to Fe III, which operates for the majority of iron atoms in line formation layers below log 5000 = -1. Unfortunately, the existing NLTE calculations for Fe are not accurate enough to be applied now to our stars. Clearly, an extended energy-level model atom is needed for a reliable non-LTE analysis of Fe. The ionisation equilibrium for different elements/ions (or its violation) can be seen in Table 4 and is discussed in Sect. 4.
3.3. Spectrophotometry

Fe abundance

excitation potential (eV)
Fig. 3. Iron abundance vs. equivalent widths (upper panel) and excitation potential (lower panel) for 21 Peg. The open circles indicate Fe I, while the open triangles indicate Fe II.

Figure 3 shows the correlations of Fe I and Fe II abundances with the equivalent widths (upper panel) and with the excitation potential (lower panel) for 21 Peg. The correlation with the equivalent widths is shown for a mic of 0.5 km s-1 , which we infer to be the best value for 21 Peg, since the slope of the linear fit for Fe I is -4.897 â 10-3 ± 1.560 â 10-3 må-1 and for Fe II is 2.146 â 10-5 ± 4.058 â 10-4 må-1 . Here we gave a preference to the result obtained from Fe II because of the higher number of measured Fe II lines in a wider range of equivalent widths. The same analysis was made for HD 145788 and for Cet. The error bar on mic was calculated using the error bar of the slope (abundance vs. equivalent width) derived from a set of different mic .

For a complete self-consistent analysis of any star, one should reproduce the observed spectral energy distribution with the adopted parameters for a model atmosphere. In the optical spectral region, spectrophotometry was only available for 21 Peg and Cet, while in the ultraviolet, IUE spectra were available for all three stars. For Cet ultraviolet spectrophotometry from the TD1 satellite (Jamar et al. 1976) was also available, along with the flux calibrated spectra from STIS at HST (Gregg et al. 2004). The comparison between the observed flux distributions and the model fluxes calculated with the adopted atmospheric parameters for 21 Peg and Cet is shown in Fig. 4. For HD 145788 we estimated a reddening E(B-V)0.2 from the dust maps of Schlegel et al. (1998). The comparison of reddened model fluxes with the available IUE spectrum, Johnson UBV photometry (Nicolet 1978), Geneva photometry3 and 2MASS photometry (Cutri et al. 2003) is shown in Fig. 11 (online material). This plots supports the value of E(B-V)=0.2 and shows good agreement between all the observations and the model fluxes, confirming the obtained fundamental parameters, and also the importance of considering reddening in the analysis of relatively nearby stars such as HD 145788. The optical spectrophotometry was taken from Adelman et al. (1989) and Breger (1976). All flux measurements were nor3

http://obswww.unige.ch/gcpd/ph13.html


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Fossati et al.: The abundance analysis for normal A- and B-type stars.

1.2 1.1 1 0.9

normalised flux

0.8

0.7 0.6 0.5 0.4 0.3 0.2

4850

4860

4870
o

wavelength (A)
malised to the flux value at 5000 å obtained from observed data for Cet and 21 Peg and from the model fluxes for HD 145788. Also for 21 Peg and Cet, we extended the comparison to the near infrared region including 2MASS photometry (Cutri et al. 2003). Following Netopil et al. (2008), we obtained a reddening E(B-V) 0.01 ± 0.01 mag for both Cet and 21 Peg. A good agreement between the unreddened model fluxes and the observed spectrophotometry from UV to near infrared confirms negligible reddening for these two stars. The agreement between the IUE observation and the model fluxes is very good for HD 145788, but the IUE data alone cannot be used as a check of T eff , since a variation of about 1000­

Fig. 2. Observed H line profile for 21 Peg (black solid line), compared to synthetic profiles (red, blue, and green lines). The red solid lines were obtained assuming the best T eff and log g values of Table 1. The blue lines show the synthetic profile by increasing T eff by 200 K (top spectrum) or log g by 0.1 dex (bottom spectrum). The green lines show the synthetic profile by decreasing T eff by 200 K (top spectrum) or log g by 0.1 dex (bottom spectrum).

1500 K is needed to make any visible discrepancy. At the T eff of HD 145788 hydrogen lines were still quite sensitive to temperature variations, so we did not need accurate spectrophotometry for a reliable effective temperature estimate. The slight disagreement between the model fluxes and the 2MASS photometry could stem from our not accounting for reddening in the calculation of the model fluxes since the derived value is too uncertain. The model fluxes of 21 Peg are in good agreement with the observations from the ultraviolet up to the near infrared. For 21 Peg and Cet, few data points in the red appear above the model fluxes. This often happens with the optical spectropho-


Fossati et al.: The abundance analysis for normal A- and B-type stars.

7

-2 Cet

-1

m-m(5000)

Cet

0

21 Peg 1 Adelman Breger TD 1 model fluxes IUE STIS 2MASS 2000 6000 10000 20000

2

wavelength in log scale (A)
tometry as shown in Fig. 1 of Adelman et al. (2002). We do not have a definite explanation for this effect, although taking a very accurate measurement of the reddening into account could remove this inconsistency. As for the hydrogen lines, we checked the fit of the spectral energy distribution to model fluxes calculated with a combination of a lower temperature and lower gravity or a higher temperature and higher gravity. In both cases the fit clearly becomes worse mainly in the spectral region around the Balmer jump. This definitively excludes the possibility that a combination of fundamental parameters, which gives a worse but still reasonable fit to the hydrogen line profiles, should considered.

o

Fig. 4. Comparison between LLMODELS theoretical fluxes calculated with the fundamental parameters and abundances derived for 21 Peg and Cet (green solid line) with spectrophotometries by Adelman et al. (1989) (open red circles) and Breger (1976) (open blue triangles), IUE calibrated fluxes (full black line), TD 1 observations (full blue squares) and STIS spectrum (full red line). The violet full dots represent the 2MASS photometry for which the given error bar in wavelength shows the wavelength range covered by each filter. The model fluxes for 21 Peg and Cet, shown in the middle of the panel, were convolved to have the same spectral resolution of the IUE fluxes (R 900), while the upper of Cet has a spectral resolution 700, approximately the same of STIS spectra. For a better visualisation of the whole spectral region the x-axis is on a logarithmic scale.

For Cet we obtained several spectrophotometric observations from different sources. In particular, the data in the optical region around the Balmer jump were conclusive for T eff determination, as mentioned already in Sect. 3.2.1. The STIS spectrum shown in Fig. 4 is formed by three separated spectra (UV, visible, IR) with about the same spectral resolution. We noticed a small (but not negligible) vertical jump at the overlapping wavelengths between the three spectra. For this reason we decided to fit them separately. This decision was based on an adjustment performed ´ at the level of the flux calibration (Ma´ iz-Apellaniz 2005, 2006). The use of an energy distribution plays a crucial role in adjusting the atmospheric parameters because it is independent of the photometric calibrations (different for each author), of the


8

Fossati et al.: The abundance analysis for normal A- and B-type stars.

hydrogen line fitting (reduction and normalisation dependent), and of the ionisation equilibria approach (dependent on several effects such as the adopted atomic line data and non-LTE effects). The uncertainties on T eff and log g were derived from the hydrogen line fitting. This way of deriving the error bars on the parameters also includes the SNR of the observations. For all three stars, we derived an error bar on the T eff of 200 K and 0.1 dex in log g, as also shown in Figs. 8, 9, and 10 (online material).
3.4. Comparison with previous determinations

tometry, essentially the same way as we did. The main difference is that they only used one available spectrophotometric set, while we used all the data found in the literature. The spectrophotometry by Breger (1976) appears a little below the one by Adelman et al. (1989) around the Balmer jump. To be able to simultaneously fit both sets of measurements, we needed a T eff lower than 13000 K and the best fit was obtained for T eff = 12800 K. Also our H profile is best-fitted only with a T eff below 13000 K, as already mentioned in Sect. 3.2.1.

Only one previous temperature determination for HD 145788 is known. Glagolevskij (1994) determined a T eff of 9600 K from the reddening free Q parameter (Johnson & Morgan 1953) and of 9100 K from the X parameter derived from multicolour photometry. The T eff derived from the Q parameter agrees with our estimation. The star 21 Peg was analysed for abundances and atmospheric parameters several times in the past. The atmospheric parameters extracted from literature are collected in Table 2, together with the methods of their determination. The last column of Table 2 lists the methods adopted to derive the fundamental parameters for each author. SPh and JPh correspond to ¨ Stromgren and Johnson photometry respectively. "Fe eq" indicates the use of the Fe ionisation equilibrium, while H-fit the use of hydrogen line fitting. SED corresponds to the use of spectral energy distribution in the visible and/or UV region. All the values obtained from the literature are in excellent agreement with our adopted parameters. We would like to mention that only the oldest determination (Sadakane 1981) was obtained using all the possible parameter indicators, as done in this work.
Table 2. Atmospheric parameters of 21 Peg derived from other authors. Reference S81 B82 A83 A83 A83 M85 S93 L98 W00 H01 A02 A02 F09 T eff [K] 10500 10500 10700 10350 10375 10300 10450 10200 10350 10450 10375 10350 10400 log g [cgs] 3.50 3.50 (...) (...) (...) (...) 3.50 3.50 3.48 3.60 3.47 3.55 3.55 method SPh, JPh, Fe eq, H-fit, SED SPh, H-fit JPh SPh SED SED SPh, H-fit SPh SPh SPh SPh SED, H-fit SED, H-fit

Table 3. Atmospheric parameters of Cet derived from other authors. Reference H79 K80 M83 M83 A84 M85 S88 M88 R90 A91 T91 S92 A02 A02 F09 T eff [K] 13100 13000 13030 13170 13150 13170 13200 13425 13150 13150 13000 13210 13174 13100 12800 log g [cgs] 3.90 (...) 3.88 3.91 3.65 (...) 3.70 (...) 3.50 3.85 (...) 3.65 3.70 3.85 3.75 method (as in Table 2) SPh SED SED(UV) SED(visible) SED, H-fit SED SED SPh SPh H-fit SED SPh SPh SED, H-fit SED, H-fit

H79: Heacox (1979), K80: Kontizas & Theodossiou (1980), M83: Malagnini et al. (1983), A84: Adelman (1984), M85: Morossi & Malagnini (1985), S88: Sadakane et al. (1988) M88: Megessier (1988), R90: Roby & Lambert (1990), A91: Adelman (1991) T91: Theodossiou & Danezis (1991), S92: Singh & Castelli (1992), A02: Adelman et al. (2002) F09: this work. The methods and the acronyms are as in Table 2.

4. Abundance analysis
The VALD database is the main source for the atomic parameters of spectral lines. For light elements C, N, O, Ne I, Mg I, Si II, Si III, S, Ar, and also for Fe III, the oscillator strengths are taken from NIST online database (Ralchenko et al. 2008). LTE abundance analysis in the atmospheres of all three stars is based mainly on the equivalent widths analysed with the improved version of WIDTH9 code updated to use the VALD output linelists and kindly provided to us by V. Tsymbal. In the case of blended lines or when the line is situated in the wings of the hydrogen lines, we performed synthetic spectrum calculations with the SYNTH3 code. For our analysis we used the maximum number of spectral lines available in the observed wavelength ranges except lines in spectral regions where the continuum normalisation was too uncertain (high orders of the Paschen series in the ESPaDOnS spectrum of Cet, for example). The final abundances are given in Table 4. A line-by-line abundance list with the equivalent width measurements, adopted oscillator strengths, and their sources is given in Table 9 (online material). Below we discuss the results obtained for individual elements.

S81: Sadakane (1981), B82: Boesgaard et al. (1982), A83: Adelman & Pyper (1983), M85: Morossi & Malagnini (1985), S93: Smith & Dworetsky (1993), L98: Landstreet (1998), W00: Wahlgren & Hubrig (2000), H01: Hubrig & Castelli (2001), A02: Adelman et al. (2002) F09: this work.

The spectroscopic literature for Cet is extremely vast and started in the early 60s. We decided to compare the determinations of T eff and log g from the late 70s on. These data and our determination are presented in Table 3. Our value for T eff is the only one below 13000 K, while log g agrees with all the other measurements. Adelman et al. (2002) also derived the parameters of Cet with the simultaneous fitting of H and spectropho-


Fossati et al.: The abundance analysis for normal A- and B-type stars.

9

Table 4. LTE atmospheric abundances in programme stars with the error estimates based on the internal scattering from the number of analysed lines, n. Ion He I CI C II NI N II OI O II Ne I Na I Mg I Mg II Al I Al II Al III Si I Si II Si III P II P III S II Cl II Ar I Ar II Ca I Ca II Sc II Ti II V II Cr I Cr II Mn I Mn II Fe I Fe II Fe III Co II Ni I Ni II Zn I Sr II Y II Zr II Ba II Nd III T eff log g HD 145788 log(N /Ntot ) -1.10±0.05 -3.60±0.14 -3.63: -3.12±0.09 21 Peg log(N /Ntot ) -1.09±0.03 -3.66±0.14 -3.65±0.05 -3.95±0.12 -3.90: -3.28±0.11 -3.76±0.05 -5.60: -4.42±0.12 -4.56±0.03 -5.89: -5.70±0.10 -4.95: -4.49±0.13 -4.26±0.18 -6.37±0.06 -4.86±0.13 Cet log(N /Ntot ) -0.97±0.04 3.58±0.07 4.03±0.13 3.74±0.07 3.06±0.14 -3.04: -3.66±0.09 -5.23±0.07 -4.27: -4.47±0.16 -5.57: -5.73±0.27 -5.30±0.02 -4.80: -4.41±0.20 -4.16: -6.38±0.19 -6.19: -4.78±0.16 -6.95: -4.86±0.24: -5.24±0.19 -5.77: -9.31: -7.42±0.08 -6.41±0.10 - - - - 6.50±0.09 4.53±0.22 4.58±0.14 4.52±0.10 -6.93: - - - - Sun (*) log(N /Ntot ) -1.12 -3.65 -3.65 -4.26 -4.26 -3.38 -3.38 -4.20 -5.87 -4.51 -4.51 -5.67 -5.67 -5.67 -4.53 -4.53 -4.53 -6.68 -6.68 -4.90 -6.54 -5.86 -5.86 -5.73 -5.73 -8.99 -7.14 -8.04 -6.40 -6.40 -6.65 -6.65 -4.59 -4.59 -4.59 -7.12 -5.81 -5.81 -7.44 -9.12 -9.83 -9.45 -9.87 -10.59 5777 K 4.44

n 4 5 2 7

n 7 9 4 4 1 18 7 1 5 7 2 4 1 22 2 3 26

n 6 7 10 9 9 2 20 3 2 10 2 8 3 1 31 2 9 1 31 2 2 6 2 1 11 21 3 7 186 4 1 17 2

-4.17±0.24 -4.29±0.06 -5.70: -5.11±0.10 -4.75: -4.27±0.14

5 4 2 3 1 11

-4.36:

2

-5.46±0.11 -5.54±0.16 -8.90±0.03 -6.80±0.15 -7.55±0.20 -6.15±0.09 -5.86±0.11 -6.43: -6.15±0.08 -4.23±0.16 -4.13±0.15 -5.32±0.13 -5.12±0.05 -8.45: -9.06: -8.75±0.32 -8.96±0.11 9750 K 3.7

5 6 5 47 4 4 31 2 4 88 147 8 3 2 1 3 3

5.84±0.11 3 5.98±0.08 5 9.37±0.10 7 7.23±0.09 59 7.98±0.06 9 6.29±0.09 5 6.20±0.10 68 6.54±0.21 5 6.51±0.17 19 4.52±0.13 108 4.50±0.12 406 4.60±0.06 3 6.75±0.18 3 5.71±0.05 10 5.61±0.09 23 -6.86: 1 -9.10: 2 -9.76±0.15 4 -9.48±0.28 4 -9.19±0.06 3 -10.09±0.07 3 10400 K 3.55

- - - - - - - - - - - - - - -

-5.76±0.19 -9.15:

12800 K 3.75

Internal scattering was not estimated when n < 3, in which case the derived abundance if flagged with a colon (:). (*) the abundances of the solar atmosphere calculated by Asplund et al. (2005).

4.1. Results for individual elements 4.1.1. Helium

He I 4472 å line is shown in Fig. 4.1.1 for different temperatures in 21 Peg (Online material). The helium abundance in HD 145788 and in 21 Peg is solar, while it is slightly overabundant in Cet. Our analysis was applied to the He I lines at wavelengths shorter than 5000 å, which should be influenced very little by non-LTE effects, except, maybe, in the case of Cet where LTE synthetic profiles fit the line wings but not the line cores. Non-LTE calculations for He I lines in the spectrum of Ori (T eff =13000 K, log g=2.0) show that negative non-LTE corrections of about 0.1­0.2 dex should even be applied to blue lines (Takeda 1994). It is unclear what corrections are expected for main sequence stars of

Stark broadening of helium lines was treated using the Barnard et al. (1974) broadening theory and tables. For allowed isolated lines we used width and shift functions from Table 1 of this paper, while an interpolation of the calculated line profiles given in Tables 2-8 was employed for 4472 å line. All abundances were derived without using equivalent widths, but by fitting of the line wings. The quality of the fit to the observed He I lines in the programme stars is demonstrated in Fig. 1, while the fit to


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Fossati et al.: The abundance analysis for normal A- and B-type stars.

the same T eff , therefore non-LTE calculations for at least Cet are necessary for deriving He abundance with proper accuracy. Our high-quality observations of Cet will serve perfectly for a thorough non-LTE study of He I lines in middle B-type stars.
4.1.2. CNO

Przybilla et al. (2008). In HD 145788 and Cet, oxygen seems to be slightly overabundant with values falling in the solar photospheric range defined by Grevesse et al. (1996) and Asplund et al. (2005). For Cet our oxygen abundance agrees perfectly with that obtained by Roby & Lambert (1990).
4.1.3. Neon and argon

The carbon abundance as derived from C II lines is solar for all three stars. It is also very close to the cosmic abundance standard recently determined by Przybilla et al. (2008) from the analysis of nearby early B-type stars and discussed in their paper. The ionisation equilibrium between C I and C II deserves some short comment. Przybilla et al. (2001a) calculated non-LTE corrections for C I and C II for Vega model atmosphere (T eff =9550 K, log g=3.95). They found these corrections to be negligible. Rentzsch-Holm (1996) made a non-LTE analysis of C II in Atype stars also obtaining very small (less than 0.05 dex) negative corrections at effective temperatures around 10000 K. Because the ionisation equilibrium between C I/C II is fulfilled for HD 145788, we assume that non-LTE corrections are also negligible for this star. Instead, at higher T eff values, the abundances obtained from C I lines having the lower level 3s1 Po ( 4932, 5052, 5380, 8335, 9406 å) are significantly lower than those obtained from the other C I lines. In the case of 21 Peg, abundances of C I and C II only agree if we neglect the abundances obtained from 4932, 5052, 5380 lines. In the spectrum of the hottest star of our sample, Cet, the situation is even more extreme. 4932, 5052, 5380 C I lines are not visible at all, while 8335, 9406 å lines appear in emission. The C I lines at 7111-7120 å are rather shallow, we can only determine an upper limit for the abundance: log(C/Ntot ) = -4.0. Like Roby & Lambert (1990) we also obtain a C I/C II imbalance in Cet. Non-LTE calculations by Rentzsch-Holm (1996) seem to explain the unusual behaviour of C I 4932, 5052, 5380 lines in stars hotter than Cet because the abundance corrections become positive and grow with the effective temperature. Nitrogen abundance is obtained for the two hottest stars of our programme from the lines of the neutral and singly-ionised nitrogen. While in 21 Peg we get the evidence for ionisation equilibrium, in Cet N II lines give higher abundance by 0.3 dex. In both stars, LTE nitrogen abundance exceeds the solar one. For Cet, Roby & Lambert (1990) derived an approximately solar nitrogen abundance from N I lines located between two Paschen lines. We use these lines, too, and the higher nitrogen abundance derived by us is caused by the larger equivalent widths, and not by the difference in the adopted effective temperature. From the non-LTE calculations performed by (Przybilla & Butler 2001), we may expect -0.3 abundance corrections for the lines of N I that bring nitrogen abundance in both stars to the solar value. It is difficult to estimate corrections for N II lines. Evidently, non-LTE analysis of both C I/C II and N I/N II line formation is necessary. Oxygen abundance was derived from the lines of neutral oxygen in HD 145788 and in 21 Peg, while the lines of neutral and singly-ionised oxygen were used in Cet. Although there are plenty of O I lines in the red region, our analysis was limited by the lines with 6500 å, which are not influenced by non-LTE effects or very little so (Przybilla et al. 2000). Even for 6155­8 å, lines the abundance corrections are less than 0.1 dex in main sequence stars. Within the errors of our abundance analysis, 21 Peg has nearly solar oxygen abundance. Moreover, it agrees perfectly with the cosmic abundance standard derived by

The abundance of these noble gases in stellar atmospheres attracts special attention because they cannot be obtained directly in the solar atmosphere. These gases are volatile, and meteoritic studies also cannot provide the actual abundance in the solar system. The revision of the solar abundances by Asplund et al. (2005) that results in a 0.2­0.3 dex decrease in CNO abundances, so those of other elements produced significant inconsistency between the predictions of the solar model and the helioseismology measurements. One of the ways to bring both data into agreement is to increase Ne abundance. Solar model calculations by Bahcall et al. (2005) show that A(Ne)=8.29±0.05 is enough for this purpose. Cuhna et al. (2006) performed non-LTE analysis of neon line formation in the young B-type stars of the Orion association and derived an average A(Ne)=8.11±0.05 (log(Ne/Ntot ) = -3.93) from 11 stars. Hempel & Holweger (2003) derived the Ne abundance in a sample of optically bright, early B-type main sequence stars, obtaining an average non-LTE Ne abundance of log(Ne/Ntot ) = -3.93 ± 0.13. Przybilla et al. (2008) obtained non-LTE Ne abundances in a sample of six nearby main sequence early B-type stars. They derived an Ne abundance of A(Ne)=8.08±0.03 (log(Ne/Ntot ) = -3.96). All these values agree very well with A(Ne)=8.08±0.10 obtained from analysing the emission registered during low-altitude impulsive flare (Feldman & Widing 1990), but are 0.3 dex higher than adopted by Asplund et al. (2005). Finally, Ne I and Ne II nonLTE analysis in 18 nearby early B-type stars (Morel & Butler 2008) results in A(Ne)=7.97±0.07. All determinations in B-type stars agree within the quoted errors and provide the reliable estimate of neon abundance in the local interstellar medium, which is higher than the newly proposed solar neon abundance. Recently, Lanz et al. (2008) have studied the Ar abundance in the same set of young B-type stars of the Orion association and derived an average A(Ar)=6.66±0.06 (log(Ar/Ntot ) = -5.38), which again agrees with the value A(Ar)=6.57±0.12 obtained by Feldman & Widing (1990), but is 0.4 dex higher than recommended by Asplund et al. (2005). While non-LTE effects on Ne I lines are known to be strong (see Sigut 1999), those on Ar II lines are weak, 0.03 dex (Lanz et al. 2008). We measured Ne I lines in the spectra of 21 Peg and Cet and Ar I/Ar II in Cet only. As expected, averaged LTE neon abundances in both stars are higher than the solar one and than that derived by Hempel & Holweger (2003), Cuhna et al. (2006), Morel & Butler (2008), or Przybilla et al. (2008) for B-type stars. However, applying non-LTE corrections, calculated by Dworetsky & Budaj (2000) for the strongest Ne I 6402 å line to our LTE abundances derived from this line in both stars we get log(N e/Ntot ) = -3.89 and -3.86 for 21 Peg and Cet, respectively, which brings Ne abundance in both stars into rather good agreement with the results obtained for early B-type stars. Rough estimates of possible non-LTE corrections in Cet for Ne I 6402 and 6506 å lines, for which LTE and non-LTE equivalent widths versus effective temperature are plotted by Sigut (1999), give us N eLT E - N eN LT E -0.3 dex for each line, and it agrees with the correction -0.36 calculated by Dworetsky &


Fossati et al.: The abundance analysis for normal A- and B-type stars.

11

Budaj (2000) for Ne I 6402 line. The star Cet is a young star and thus adds reliable current data on the Ne abundance, taking a large number of high-quality line profiles and secure model atmosphere parameters into account. We derived argon abundance log(Ar/Ntot ) = -5.24 ± 0.19 in Cet from 5 weak but accurately measured Ar II lines. Within error bars this value agrees with the results by Lanz et al. (2008) for B-type stars in the Orion association. We also managed to measure the two strongest Ar I lines at 8103 and 8115 å. They each give an Ar abundance that is too high. The 8115 å line is blended with a Mg II line, and taking this blend into account we get log(Ar/Ntot ) = -5.2, while the 8103 å line is too strong for its transition probability. Moreover, both lines are located in the region contaminated by weak telluric lines, therefore the extracted abundances may be uncertain.
4.1.4. Na, Mg, Al, Si, P, S, Cl

in the measurements is of the order of the cited accuracies, and theoretical calculations agree rather well with the experimentally measured transition probabilities and Stark widths. Phosphorus is overabundant relative to the solar abundance by 0.3 dex in both hotter stars, 21 Peg and in Cet. Oscillator strengths for P II lines (taken from VALD) originally come from calculations by Hibbert (1988). Therefore, at least part of the observed overabundance may be caused by uncertainties in calculated transition probabilities. No non-LTE analysis is available for the phosphorus line formation. At the limit of detection, we managed to measure the two strongest Cl II lines (at 4794, 4810 å) in Cet. The obtained upper limit on the chlorine abundance is 0.4 dex lower than the recommended solar value (Asplund et al. 2005). Sulphur is overabundant by 0.5 dex in HD 145788 and almost solar in the two other stars.
4.1.5. Ca and Sc

In both 21 Peg and Cet, Na abundances are derived from lines affected by non-LTE(Takeda 2008), therefore it is not surprising that their measured values are different from the solar ones. In all three stars we obtained a slight Mg I/Mg II imbalance. For 21 Peg and Cet, the abundance derived from Mg II lines is close to solar. The abundance derived from the weaker Mg I lines is consistent to the one obtained from Mg II, hence solar. We conclude that, for these two stars, Mg abundance is consistent with the solar one and that discrepancies observed in the strong Mg I lines come from non-LTE effects. Magnesium is slightly overabundant in HD 145788, but all the Mg I lines are affected by non-LTE effects. The sign and the magnitude of its effect depends on both T eff and log g (Przybilla et al. 2001b), which prevents us from obtaining any firm conclusion, until detailed non-LTE calculations for Mg I for early Aand middle B-type main sequence stars are carried out. Aluminum is above solar in HD 145788 and have nearly solar abundances in the two other stars as derived from Al II lines. Al I lines in all stars and Al III lines in Cet provide some discordant results. It is not possible to discuss the ionisation equilibrium without a non-LTE analysis of the line formation of all three ions. Accurate silicon abundance determination in stars and the interstellar medium is an important part of abundance studies and intercomparisons because Si is a reference element for meteoritic abundances. Silicon is slightly above solar in HD 145788, and close to solar in 21 Peg and in Cet, if we consider the results obtained from the numerous Si II lines. The spectral synthesis in the region of the only Si I 3905 å line observed in all three stars provides much lower Si abundance, and the abundance difference between Si I and Si II is practically independent of T eff . A non-LTE analysis of Si line formation in the Sun and in Vega (Wedemeyer 2001) shows that positive non-LTE corrections are expected for Si I 3905 å line, while small negative corrections may be expected for Si II lines, thus leading both ions into the equilibrium. Our Si analysis is based on the very accurate transition probability for Si I 3905 å line (O'Brian & Lawler 1991) and on a combination of transition probabilities extracted from a recent NIST compilation (Kelleher & Podobedova 2008) and theoretical calculations by Artru et al. (1981) for Si II lines. The NIST compilation does not contain data for about a quarter of the lines observed in our stars for which rather concordant data exist. Table 6 (online material) gives a collection of the experimental, as well as theoretical, atomic parameters for Si II lines that may be useful in a future non-LTE analysis. The dispersion

These two elements are of special interest in A-type star studies, as their non-solar abundances indicate of a star's classification as a metallic-line (Am) star. In hot Am stars with T eff close to HD 145788, both elements, in particular scandium, are underabundant by 0.4­0.5 dex, while other Fe-peak elements are overabundant by 0.2­0.3 dex (see for instance o Peg, which is a typical representative of the hot Am stars, Adelman 1988). Calcium and scandium are overabundant in HD 145788, but underabundant in 21 Peg by 0.2 dex and 0.4 dex, respectively. Formal Ca-Sc classification criteria of Am stars would make 21 Peg the hottest known Am star. However, classical Am stars are also characterised by overabundances of 0.2 - 0.3 dex for other Fe-peak elements, and even more remarkable overabundances of Sr, Y, and Zr (Fossati et al. 2007). All these peculiarities are not observed in 21 Peg, which therefore cannot be classified as Am. The star Cet has solar Ca abundance and the same Sc deficiency as 21 Peg. At the T eff of 21 Peg non-LTE corrections for both Ca I and Ca II, lines are expected to be positive (L. Mashonkina, private communication). In other words, the Ca abundance obtained from LTE calculations is, perhaps, underestimated, which may explain the observed discrepancies with respect to the solar case. Detailed non-LTE analysis of the formation of Ca lines is required for accurate abundances. Including hyperfine splitting (hfs) does not change abundance results because hfs-effects are negligible for the investigated Sc II lines (see Kurucz' hfs calculations 4 ). Scandium deficiency requires more careful non-LTE analysis, because it is observed not only in classical Am stars, but also in other stars, for example, in the A-type supergiant Deneb (Schiller & Przybilla 2008), which have solar abundances of the other Fe-peak elements.
4.1.6. Ti, V, Mn, Co, Ni

Within the error limits, all these elements are almost solar in 21 Peg. The same is true for Cet, except for Ti, which is slightly underabundant. The Ti abundance determinations are based on the accurate laboratory transition probabilities (Pickering et al. 2001) currently included in VALD. The non-LTE corrections are expected to be positive (Schiller & Przybilla 2008), leading to abundance values closer to the solar one. The situation is a bit different in the atmosphere of HD 145788, where all these elements exhibit 0.2­0.4 dex over4

http://cfaku5.cfa.harvard.edu/ATOMS


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Fossati et al.: The abundance analysis for normal A- and B-type stars.

abundance relative to the solar photospheric abundances. (This would indicate that the star is Am, but since Ca and Sc are not underabundant, the star cannot be classified as Am). Still in HD 145788, for Mn, Ni, Cr, and Fe, the lines of the first ions provide slightly higher abundance than the lines of the neutrals, while no significant difference in ionisation equilibrium is observed in the two other stars. The Mn lines are known to have rather large hfs. We checked the influence of hfs on the derived Mn abundances for HD 145788 where we measured the largest equivalent widths. Data on hfs for Mn were taken from Blackwell-Whitehead et al. (2005) (Mn I) and from Holt et al. (1999) (Mn II). We found that this effect is weak, and does not exceed 0.05 dex even for the lines with the largest hfs.
4.1.7. Cr

Although based only on a few lines, abundances derived from Cr I lines are accurate because the recommended atomic parameters of these transitions (Martin et al. 1988) currently included in VALD are supported by recent precise laboratory measurements (Sobeck et al. 2007). For Cr II, laboratory measurements are only available for the low-lying lines with E i 4.8 eV. Few measured lines in HD 148788 and in Cet, and about one third of Cr II lines in 21 Peg, originated in levels with higher excitation potential. For these lines, only theoretical calculations are available. In several publications, favour was given to the transition probability calculations performed with the orthogonal operator technique (RU: Raassen & Uylings 1998), which are collected in the RU database5 for the Cr II, Fe II, and Co II ions. All these data are included in the current version of VALD. A comparison between RU calculations and the most recent laboratory analysis of 119 lines of Cr II in 2055­4850 å spectral region (Nilsson et al. 2006) shows that both sets agree within 10% on the absolute scale with a dispersion of 0.13 dex. In the optical spectral region, Nilsson et al. (2006) measures only 7 lines (in 4550­4850 å region). To provide a consistent analysis, we decided to use RU data for all lines of Cr II in our work. Ionisation equilibrium was found for 21 Peg, while a slight imbalance was found between Cr I and Cr II in HD 145788. In the spectrum of Cet no Cr I lines are present. We found that Cr is overabundant in HD 145778 by 0.4 dex (as average between the Cr I and Cr II abundances), slightly overabundant in 21 Peg (by 0.15 dex), while it has solar abundance in Cet.
4.1.8. Fe

21 Peg spectrum, and more than 200 lines have excitation potential > 10 eV. In Cet, 98 out of 186 measured Fe II lines have excitation potential > 10 eV. For all these high-excitation lines, atomic parameters are available only through theoretical calculations. As for Cr II lines, we checked the RU database. For 66 lines for which laboratory transition probabilities are measured, a comparison with the RU data results in logg f (lab data) - logg f (RU data) = 0.11 ± 0.11 dex. This means that using the Fe II lines of the homogeneous set of transition probabilities obtained from theoretical calculations and available in the RU database, we may over or underestimate the corresponding abundances by no more than 0.1 dex. Only calculated transition probabilities are available for Fe III lines. For a comparison of the accuracy of Fe II lines atomic parameters, we obtained three sets of abundance determinations for 21 Peg: one based on laboratory data included in VALD, one based on the recent NIST compilation (Ralchenko et al. 2008), and one based on RU data. The results are (i) from laboratory VALD data: log(Fe ii/Ntot ) = -4.61 ± 0.11 (51 lines); (ii) from NIST data: log(Fe ii/Ntot ) = -4.52 ± 0.17 (68 lines); (iii) from the same set of RU data: log(Fe ii/Ntot ) = -4.49 ± 0.09 (51 lines). Few strong high-excitation Fe II lines are included in the NIST compilation. These results justify the use of RU calculations for accurate iron abundance analysis, when laboratory measurements are not available. High accuracy of spectral data, together with very low sin i, and fairly well-established atmospheric parameters, makes 21 Peg a perfect object for an Fe II study. We could measure practically all unclassified lines with intensity 1 and higher given in the list of laboratory measurements (Johansson 1978). These lines belong to the transitions with very high excitation potentials. Accurate position and intensity measurements in stellar spectra may help in further studies of the Fe II spectrum and term system. Discussion of the possible non-LTE effects on Fe I and Fe II lines was given at the end of Sect. 3.2.2. LTE iron abundances in 21 Peg and in Cet agree well with the cosmic abundance standard (see Przybilla et al. 2008).
4.1.9. Sr, Y, Zr

These elements are overabundant in HD 145788 and have solar abundances in 21 Peg and Cet. For the Zr analysis we used the most recent experimental transition probabilities from Ljung et al. (2006).
4.1.10. Ba and Nd

In 21 Peg and in Cet, iron abundance is practically solar. In particular, the same Fe abundance is derived from Fe lines in three ionisation stage for Cet. No obvious ion imbalance was detected. Iron is a crucial element for adjusting microturbulence, metallicity, and model atmosphere parameters. It has the most spectral lines in the first three ionisation stages with relatively accurate atomic data that can be observed in the optical spectra of early A- and middle B-type stars. Numerous Fe II lines in the range of excitation energy 2.5­11.3 eV are seen in the spectrum of 21 Peg. In the 4000­8000 å spectral region, laboratory measurements are available for 66 Fe II lines with Ei 6.2 eV (Ryabchikova et al. 1999). We measured 406 Fe II lines in the
5

ftp://ftp.wins.uva.nl/pub/orth

Barium is overabundant in HD 145788 and in 21 Peg. No lines of elements heavier than zirconium are identified in our hottest programme star Cet. While barium overabundance, together with strontiumyttrium-zirconium overabundances in HD 145788, favours its classification as a hot Am star, barium overabundance in 21 Peg, which otherwise has near solar abundances of practically all other elements, is unexpected. Non-LTE corrections to barium abundance, if any, should be positive (L. Mashonkina, private communication). Gigas (1988) derived for Vega non-LTE corrections for the two barium lines ( 4554, 4934 å) measured in this work as well. They obtained a positive correction of about 0.3 dex for both lines. The non-LTE corrections that should be applied to the barium abundance obtained in HD 145788 and in 21 Peg make the problem of the barium overabundance even


Fossati et al.: The abundance analysis for normal A- and B-type stars.

13

more puzzling. Practically in all the abundance studies of normal A-type stars Ba was found to be overabundant (Lemke 1990; Hill & Landstreet 1993). We measured three weak features at the position of the strongest Nd III lines (Ryabchikova et al. 2006) that result in slight Nd overabundance in 21 Peg. While Nd overabundance may still be attributed to the uncertainties on the absolute scale for calculated transition probabilities or non-LTE effects, which are expected to be negative (Mashonkina et al. 2005), it is not the case for Ba where the atomic parameters of the lines using in our analysis are accurately known from laboratory studies.
4.2. Abundance uncer tainties

0.28 0.24

HD 145788 21 Peg Cet

abundance error

0.2

0.16 0.12 0.08

0.04

The abundance uncertainties for each ion shown in Table 4 are the standard deviation of the mean abundance obtained from the individual line abundances. Since our derivation of the abundances is mainly based on equivalent widths, we first have to estimate equivalent width errors given a certain SNR and sin i. With a two error bar, we derived 1.2 må for HD 145788, 0.2 må for 21 Peg and 0.5 må for Cet. These values were derived by assuming a triangular line with a depth (height of the triangulum) equal to 2 (SNR) and a width (base of the triangulum) equal to 2â sin i. The rather high uncertainty on the equivalent widths of HD 145788 is mainly due to the low SNR of its spectrum. This shows the importance of a very high SNR not only for fast rotating stars, but also for slowly rotating stars. This uncertainty includes the uncertainty due to the continuum normalisation. In Fig. 5 we plotted the error bars in abundance for a given line as a function of equivalent widths, for all the three stars analysed in this work. To derive the uncertainty in abundance due to the error bar on equivalent widths measurement, we took a representative Fe II line and derived the abundance of this line on the basis of different values of equivalent widths ranging from 0.3 må to 110 må. We then calculated the difference between the abundance obtained with the equivalent width X and X+X, where X is the error bar on the equivalent widths measurement. As expected, HD 145788 is the star that shows the largest error bar. For equivalent widths greater than 30 må, the abundance error tends asymptotically to zero. Table 9 also shows that all the lines measured for this work are above the detection limit given by the sin i and the SNR. The mean equivalent width measured in these three stars is about 20 må for HD 145788, 5 må for 21 Peg, and 8 må for Cet. These values correspond to an error bar in abundance, because of the uncertainty on the equivalent widths measurement and continuum normalisation of 0.04 dex for HD 145788, and 0.03 dex for both 21 Peg and Cet. When we have enough measured lines, we assume that the internal scatter for each ion also takes the errors due to equivalent widths measurement and continuum normalisation into account. Figure 6 shows the abundance scatter as a function of the number of measured lines for 21 Peg. For elements where nonLTE effects are supposed to be important and line-dependent, such as Al and S, the internal standard deviation is particularly high. The same is found for elements with lower accuracy in logg f values due to the complexity of the atomic levels, such as Si. For other elements with a large enough number of spectral lines, say n > 10, it is reasonable to expect an internal error of 0.11 dex (see Fig. 6).

2

4

6

8

10

12

14

16

18

20

22
o

24

26

28

30

equivalent width (mA)
Fig. 5. Error bar in abundance as a function of equivalent widths for HD 145788 (open circle), 21 Peg (open square), and Cet (open triangle). This uncertainty stems from the uncertainty in the equivalentwidth measurement and continuum normalisation. The difference in uncertainty between 21 Peg and Cet comes almost only from the difference in sin i, while the high error obtained for HD 145788 mainly from the low SNR of our observations, relative to one of the other two stars.

Considering the errors in oscillator strengths determination (see Table 9) given for the laboratory data, we may say that these errors are smaller for most elements than the internal scatter, so do not significantly influence the final results. The same is relevant for theoretical calculations. As already shown for Cr II and Fe II (Sect. 4.1.8), calculated and laboratory sets of oscillator strengths agree within 0.1 dex. It should be noted that error due to the internal scatter is just a part of the total error bar on the abundance determination. To derive a more realistic abundance uncertainty we also have to take the error bar due to systematic uncertainties in the fundamental parameters into account. For a detailed discussion of the abun-

0.21

standard deviation (dex)

0.18

0.15

0.12

0.09

0.06

0.03

10

100

number of lines
Fig. 6. Standard deviation of the derived abundances as a function of the the number of lines (shown in logarithmic scale and for a number of lines greater than 2). For visualisation reasons we omitted the standard deviation given by Zr II.


14

Fossati et al.: The abundance analysis for normal A- and B-type stars.

Table 5. Error sources for the abundances of the chemical elements of 21 Peg. Ion He I CI C II NI N II OI Ne I Na I Mg I Mg II Al I Al II Si I Si II Si III P II S II Ca I Ca II Sc II Ti II V II Cr I Cr II Mn I Mn II Fe I Fe II Fe III Co II Ni I Ni II Zn I Sr II Y II Zr II Ba II Nd III abundance log(N /Ntot ) -1.11 -3.66 -3.65 -3.95 -3.90: -3.28 -3.76 -5.60 -4.42 -4.56 -5.89 -5.68 -4.95 -4.49 -4.26 -6.37 -4.86 -5.84 -5.98 -9.37 -7.23 -7.98 -6.29 -6.20 -6.54 -6.51 -4.52 -4.50 -4.60 -6.75 -5.71 -5.61 -6.86 -9.10 -9.76 -9.48 -9.19 -10.09 abn (scatt.) (dex) 0.04 0.14 0.05 0.12 0.11 0.05 0.12 0.03 0.10 0.13 0.18 0.06 0.13 0.11 0.08 0.10 0.09 0.06 0.09 0.10 0.21 0.17 0.13 0.12 0.06 0.18 0.05 0.09 0.01 0.15 0.28 0.06 0.07 abn (T eff ) (dex) -0.08 -0.05 -0.17 -0.05 -0.09 0.00 -0.18 0.04 0.10 -0.02 0.10 -0.09 0.10 -0.07 -0.16 -0.08 -0.14 0.20 0.03 0.12 0.08 0.06 0.13 0.03 0.14 0.01 0.11 -0.02 -0.10 0.01 0.09 -0.04 0.10 0.13 0.13 0.11 0.14 0.04 abn (log g) (dex) 0.03 -0.13 -0.02 -0.07 0.05 -0.02 -0.03 -0.04 -0.06 -0.01 -0.03 0.01 -0.02 0.01 0.04 0.02 0.00 -0.07 -0.05 -0.01 0.00 0.02 -0.03 0.02 -0.03 0.01 -0.03 0.02 0.06 0.03 -0.03 0.03 -0.02 -0.03 -0.02 0.00 -0.02 0.04 abn (mic ) (dex) 0.00 -0.10 -0.08 -0.05 0.00 -0.02 -0.06 -0.01 -0.05 -0.02 0.00 -0.03 0.00 -0.03 -0.03 -0.02 -0.05 0.00 -0.04 0.00 -0.02 0.00 0.00 -0.01 0.00 -0.01 0.00 -0.02 -0.01 0.00 -0.01 -0.01 0.00 -0.06 -0.01 0.00 -0.01 0.00 abn (syst.) (dex) 0.09 0.17 0.19 0.10 0.10 0.03 0.19 0.06 0.13 0.03 0.10 0.10 0.10 0.08 0.17 0.08 0.15 0.21 0.07 0.12 0.08 0.06 0.13 0.04 0.14 0.01 0.11 0.03 0.12 0.03 0.10 0.05 0.10 0.15 0.13 0.11 0.14 0.06

Column 3 standard deviation abn (scatt.) of the mean abundance obtained from different spectral lines (internal scattering); a blank means that the number of spectral lines is < 3, hence no internal scattering could be estimated. (Note that these values are identical to those given in Table 4.) Columns 4, 5, and 6 give the variation in abundance estimated by increasing T eff by 200 K, log g by 0.1 dex, and mic by 0.4 km s-1 , respectively. Column 7 gives the the mean error calculated applying the standard error propagation theory on the systematic uncertainties given in columns 4, 5 and 6, i.e., 2 (syst.) = 2 (T eff ) + 2 (log g) + 2 (mic ). abn abn abn abn

dance error bars we again use 21 Peg. Table 5 shows the variation in abundance for each analysed ion, caused by the change of one fundamental parameter by +1, keeping fixed the other parameters. For the comparison 2 (syst.) with abn (scatt.) of those abn ions for which the internal scattering could not be measured, an a priori scatter of abn (scatt.) = 0.11 dex has to be assumed. The main source of uncertainty is the error in the effective temperature determination, while the variation due to log g and, in particular, to mic is almost negligible, although here we are considering an uncertainty on mic of 2. The abundance variation due to an increase in log g leads to a worse ionisation equilibrium for many elements for which the equilibrium is reached at the adopted log g, such as Cr, Mn, Fe, and Ni. We repeated the same analysis for Fe II for HD 145788 and Cet. The results are comparable with those obtained for 21 Peg. We also expect the same effect for the other ions. Assuming the different errors in the abundance determination are independent, we derived the final error bar using stan-

dard error propagation theory, given in column six of Table 5. When the abundance is given by a single line, we assumed an internal error of 0.11 dex. Using the propagation theory we considered the situation where the determination of each fundamental parameter is an independent process. The mean value of the LTE uncertainties given in column six of Table 5 is 0.16 dex. Because of the high SNR of the observations, the low sin i and the non peculiarity of the programme stars, we can reasonably believe that the errors in abundance determinations estimated in the present study are the smallest ones that could be obtained with the current state of the art of spectral LTE analysis for early-type stars. In the cases of stars with sin i higher than those considered in this work, the uncertainty on the abundances increases. A higher rotational velocity would force the abundance analysis to be based on strong and saturated lines that are more sensitive to mic variations than weak lines. For a more detailed discussion, see Fossati et al. (2008).


Fossati et al.: The abundance analysis for normal A- and B-type stars.

15

4.3. Comparison with previous abundance determinations

Table 6 (online) collects all previous massive abundance determinations in the atmosphere of 21 Peg (Sadakane 1981; Smith & Dworetsky 1993; Smith 1993, 1994; Dworetsky & Budaj 2000) in comparison with the results of the current analysis. We do not include works that only give abundances for a Cr and/or Fe. The main advantage of our analysis over the previous ones is the wide wavelength coverage and the high quality of our spectra. It allows us to use many more spectral lines including very weak ones of the species not analysed before. Our analysis provides homogeneous abundance data for 38 ions of 26 chemical elements from He to Nd. Reasonably good agreement exists between our abundances derived with the spectra in optical and IR spectral regions and those derived with the UV observations (Smith & Dworetsky 1993; Smith 1993, 1994), supporting the correctness of the adopted model atmosphere. Dworetsky & Budaj (2000) derived the LTE neon abundance and gave non-LTE abundance corrections for the strongest Ne I 6402 line. If we apply the non-LTE correction described in Sect. 4.1.3 to this line, we almost get the same non-LTE abundance. For Cet the number of abundance determinations in the literature is particularly vast. As for 21 Peg we consider only those where abundances are derived for large number of ions: Adelman (1991), Smith & Dworetsky (1993), Smith (1993), Smith (1994), and Acke & Waelkens (2004), except for neon abundance where we again included LTE and non-LTE results by Dworetsky & Budaj (2000). For the comparison (see the online Table 7) of the abundances obtained for Cet, we also show the adopted mic because we believe that differences in this parameter are responsible for the difference between our abundances and those of Acke & Waelkens (2004). Again, we emphasise that the total number of ions (36) and elements (22), as well as the number of lines per ion analysed in the present work, is much higher than in any previous study. For Al, Si, Fe, we managed to derive abundances from the lines of the element in three ionisation stages, which provides a unique possibility to study non-LTE effects. For the ions having several lines (S II, Ti II, Cr II and Fe II) we generally get a good agreement among the different authors, in particular, for Fe II. The lower abundances given by Acke & Waelkens (2004) are due to the high mic adopted. The difference in He abundance between our work and that by Adelman (1991) may be explained by the difference in T eff , already discussed in Sect. 3.4. Both LTE and non-LTE neon abundances agree rather well with the results by Dworetsky & Budaj (2000) after applying the non-LTE corrections given by these authors.

the possible Am classification of HD 145788, we can compare the obtained abundance pattern with the typical one for Am stars in clusters having enhanced metallicity, e.g. the Praesepe open cluster with an overall metallicity of [Fe/H]=0.14 dex (Chen et al. 2003). Fossati et al. (2007) give the abundance pattern of eight Am stars belonging to the Praesepe cluster. All these stars show clear Am signatures: underabundances of CNO and Sc, and overabundances of the elements heavier than Ti. In HD 145788 the underabundances of CNO and Sc are not visible. For this reason we believe that the star cannot be classified as a hot Am star and that the observed abundance pattern stems from the composition of the cloud where HD 145788 was formed. Table 4 and Fig. 7 show that with the adopted fundamental parameters of HD 145788 we get ionisation imbalance for the Fe-peak elements. With a small adjustment of the parameters within the error bars (T eff : +100 K, log g: -0.05) it is possible to compensate for this imbalance for the Fe-peak elements, but the ionisation balance becomes worse for other elements such as C, Mg, and Ca. This is a clear sign that a non-LTE analysis is needed for this object to understand whether the ionisation violation stems from non-LTE effects or to some other physical effect. Unfortunately we cannot be completely certain of the adopted parameters because of the lack of spectrophotometric measurements.
5.2. 21 Peg

The star 21 Peg shows solar abundances for almost all elements. At the temperature of 21 Peg, slowly rotating stars are usually hot Am, cool HgMn, or magnetic stars. The possibility of classifying 21 Peg as an Am star (which is potentially suggested by the observed Sc underabundance) is excluded by the solar abundances of all the other mentioned indicators. As explained in Sect. 4.1.5, the non-LTE correction for Ca is expected to be positive, increasing the Ca abundance to the solar value, but detailed non-LTE analysis should be performed for a more accurate determination. Almost nothing is known about non-LTE effects for Sc at these temperatures, so that to understand whether the observed underabundance is real, a non-LTE analysis should be performed. Caliskan & Adelman (1997), Adelman (1998), Adelman (1999), and Kocer et al. (2003) have published abundances of several chemically normal, early-type stars. The derived Sc abundance for stars with an effective temperature similar to that of 21 Peg is almost always below the solar value, with nearly solar abundances for other Fe-peak elements. We thus conclude that the observed abundances of most elements but Ba in 21 Peg allow us to classify it as a normal early-type rather than Am star. Solar Mn abundance and the absence of Hg II 3984 å line exclude any classification of 21 Peg as an HgMn star.
5.3. Cet

5. Discussion
In Fig. 7 we show the derived abundances normalised to solar values (Asplund et al. 2005). In the following, we discuss the stars of our sample individually.
5.1. HD 145788

HD 145788 shows a slight overabundance for almost all ions and typical Am abundance pattern for elements heavier than Ti. The overall overabundance and the Am abundance pattern could be explained if HD 145788 was formed in a region of the sky with a metallicity higher than the solar region. To be able to check

The star Cet shows solar abundances for almost all elements. Only O, Ne, Na and Ar are clearly overabundant. For Ne or Na this most probably comes from non-LTE effects. For argon, the non-LTE effects are expected to be weak (see Sect. 4.1.3). In other chemically normal, early-type stars the Ar abundance appears to be above the solar one (see Lanz et al. 2008; Fossati et al. 2007; Adelman 1998), leading to the conclusion that indirect solar Ar abundance given by Asplund et al. (2005) is underestimated.


16

Fossati et al.: The abundance analysis for normal A- and B-type stars.

1.2

HD 145788

abundance relative to the Sun

0.8 0.4 0 1.2 0.8 0.4 0 1.2 0.8 0.4 0 neutral element single ionised element double ionised element
Fig. 7. LTE abundances relative to the Sun (Asplund et al. 2005) for HD 145788, 21 Peg, and Cet, from top to bottom. The open circles, squares, and triangles indicate the abundance for the neutral element, single ionised, and double ionised, respectively. The possible nonLTE corrections for each ion are described in the text from Secs. 4.1.1 to 4.1.10.

-0.4 21 Peg

-0.4 Cet

-0.4
He C N O Ne Na Mg Al Si P S Cl Ar Ca Sc Ti V Cr Mn Fe Co Ni Zn Sr Y Zr Ba Nd

element
As for 21 Peg, the Sc underabundance of Cet is not an indication of a possible Am peculiarity. The absence of magnetic field or of normal Mn abundance, including no trace of any Hg signature in the spectrum exclude a classification of the star as magnetic peculiar or non-magnetic HgMn object. The star Cet is a known binary with a period of about 7.5 years (Lacy et al. 1997) and is a Herbig AeBe star (Malfait et al. 1998). The Herbig classification comes from a detected infrared excess at wavelengths longer than 10 µm. The two spectra obtained with ESPaDOnS show variability in the line profiles, small emission-like features close to the core of H, and emission features at the position of C I 8335 and 9405 å in the near infrared. The width of C I emission lines are exactly the same as expected for absorption line in Cet. The pre-mainsequence status of this star, which is very likely responsible for these emissions, might also explain the variability observed in the spectral lines, as circumstellar absorption or emission coming from a proto-planetary disk. The variation in the line profile within one day excludes the possibility that the observed changes are caused by the companion. Another explanation for the line variations comes from pulsation. We performed a frequency analysis of the radial velocity measurements given by Lacy et al. (1997) and the ones obtained from the two ESPaDOnS spectra. Preliminary results show that two frequencies appear in the amplitude spectrum: one corresponding to the orbital period and another one at 2.79 day-1 . This frequency is consistent with the expected pulsation periods for SPB stars with effective temperature of Cet (see Fig. 5 of Pamyatnykh 1999). In this way pulsation could explain the lineprofile variation, while the presence of a disk around the star could explain the small emission visible in H. The C I emission lines in the near infrared could be explained either by the disk or by non-LTE effects (Nieva & Przybilla 2008). Only future photometric observations and time-resolved spectroscopy and non-LTE analysis could lead to a better understanding of the star's status.

6. Conclusions: are solar abundances also a reference for early A- and late B-type stars?
One of the main goal of this work was to check whether the solar abundances can be taken as a reference for early-type stars. The same question has recently been discussed by Przybilla et al. (2008) who analysed a sample of early B-type stars in the solar neighbourhood to compare the obtained non-LTE abundances with the ones published by other authors for stars in the Orion nebula, various B-type stars, young F- and G-type stars, the interstellar medium, and the sun (Grevesse et al. 1996; Asplund et al. 2005), and to check the chemical homogeneity of the solar neighbourhood. They obtained an excellent agreement between the non-LTE abundances for He, C, N, Mg, Si, and Fe with the solar ones published by Asplund et al. (2005), while the oxygen abundance lies between the solar values obtained by Grevesse et al. (1996) and Asplund et al. (2005), and the Ne abundance is compatible with the one provided by Grevesse et al. (1996). The optical spectra of early B-type stars cannot provide reliable data for many other elements (Ca, Ti, Cr, Mn, Sr, Y, Zr), which are important for comparative abundance studies of chemically peculiar stars. While early A- and late B-type stars, investigated in the present paper, provide us with the abundances of up to 26 elements, most of which are based on enough spectral lines with accurately known atomic parameters. Figure 7 shows almost solar abundances for many elements in both 21 Peg and Cet, while the observed abundance pattern in HD 145788 gives a hint that the star may have been formed in a region of the sky at high metallicity. In early-type stars it is possible to directly derive the He abundance, while for the Sun it is only possible through astroseismological observations and modelling. For this reason it is important to check that the He abundance is comparable to the solar value in several chemically normal early-type stars. In the analysed stars, several elements (He, C, Al, S, V, Cr, Mn, Fe, Ni, Sr, Y, Zr) show abundances compatible with the revised solar data (Asplund et al. 2005), and when discrepancies


Fossati et al.: The abundance analysis for normal A- and B-type stars.

17

are present they could be explained by non-LTE effects (N, Na, Mg, Si, Ca, Ti, Nd). For Ne and Ar the expected non-LTE corrections would lead to abundances close to those derived for early B-type stars (Lanz et al. 2008; Przybilla et al. 2008) or to the solar ones given by Grevesse et al. (1996) instead of Asplund et al. (2005). Non-LTE corrections were never calculated and should be determined for other elements that show differences with the solar abundance (P, Cl, Sc, Co). We found actual discrepancies with the solar abundance for oxygen in Cet and for Ba in 21 Peg. While the oxygen problem may be solved by careful non-LTE analysis of all the available lines including the red and IR ones, the Ba overabundance cannot be explained by the current non-LTE results. The abundances obtained in this work for this set of three early B-type stars agree very well with the ones obtained by Przybilla et al. (2008) for all the elements. The published abundances of Ba in chemically normal earlytype stars (Lemke 1990; Caliskan & Adelman 1997; Adelman 1999; Kocer et al. 2003) show a definite trend towards a Ba overabundance. The non-LTE corrections for Ba should be positive, leading to an even greater discrepancy with the solar value, that probably does not represent early-type stars. Non-LTE effects are studied mainly in solar-type stars, lowmetallicity stars, and giants, and in stars hotter than early B-type, where the effects are expected to be strong. Very few analyses have been performed for normal early A- and late B-type stars (e.g. Vega), and our study claims the real need of such analyses for many elements before making a definite conclusion about the solar abundances as standards for early-type stars.
Acknowledgements. This work is based on observations collected at the Nordic Optical Telescope (NOT) as part of programme number 35-001 and at the ESO 3.6 m telescope at Cerro La Silla (Chile). Part of this work is based on observations made with the Nordic Optical Telescope, operated on the island of La Palma jointly by Denmark, Finland, Iceland, Norway and Sweden, in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias. This work is also based on observations obtained at the Canada-France-Hawaii Telescope (CFHT), which is operated by the National Research Council of Canada, the Institut National des Sciences de l'Univers of the Centre National de la Rechereche Scientifique of France, and the University of Hawaii. This work is supported by the Austrian Science Foundation (FWF project P17890-N2 - LF, TR and OK), by the Russian Foundation for Basic research (grant 08-02-00469a - TR) and by the Presidium RAS Programme "Origin and evolution of stars and galaxies" (TR). GAW acknowledges support from the Academic Research Programme (ARP) of the Department of National Defence (Canada). We thank D. Lyashko for having developed and provided the reduction pipeline for the FIES data, V. Tsymbal for providing us with an improved version of WIDTH9, and L. Mashonkina for providing some non-LTE estimates. We thank the anonymous referee for the constructive comments. We thank A. Ederoclite and L. Monaco for the spectrum of HD 145788, and M. Gruberbauer for the frequency analysis. TR and LF thank D. Shulyak for the fruitful help, support and discussion of model atmospheres and spectral energy distribution. This work made use of the MAST-IUE archive (http://archive.stsci.edu/iue/), of SAO/NASA ADS, SIMBAD, VIZIER and of the VOSpec tool (http://www.euro-vo.org/pub/fc/software.html) developed for the European Virtual Observatory. This publication makes use of data products from the Two Micron All Sky Survey, which 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.

References
Acke, B. & Waelkens, C. A&A, 2004, 427, 1009 Adelman, S. J. & Pyper, D. M. 1983, ApJ, 266, 732 Adelman, S. J. 1984, A&AS, 55, 479 Adelman, S. J. 1988, MNRAS, 230, 671 Adelman, S. J., Pyper, D. M., Shore, S. N., White, R. E. & Warren, Jr., W. H. 1989, A&AS, 81, 221 Adelman, S. J. & Gulliver, A. F. 1990, ApJ, 348, 712 Adelman, S. J. 1991, MNRAS, 252, 116

Adelman, S. J. 1998, MNRAS, 296, 856 Adelman, S. J. 1999, MNRAS, 310, 146 Adelman, S. J., Pintado, O. I., Nieva, M. F., Rayle, K. E. & Sanders, Jr., S. E. 2002, A&A, 392, 1031 Adelman, S.J., & Unsuree, N. 2007, Balt. Astron., 16, 183 Artru, M.-C., Jamar, C., Petrini, D., & Praderie, F. 1981, A&A, 96, 380 Asplund, M., Grevesse, N. & Sauval, A. J. 2005, Astronomical Society of the Pacific Conference Series, 336, 25 Auer, L. H., Mihalas, D., Aller, L. H. & Ross, J. E. 1966, ApJ, 145, 153 Bagnulo, S., Wade, G.A., Donati, J.-F., et al. 2001, A&A, 369, 889 Bahcall, J.N., Basu, S., & Serenelli, A.M. 2005, ApJ, 631, 1281 Barach, J.P. 1970, JQSRT, 10, 519 Bard, A., Kock, A. & Kock, M. 1991, A&A, 248, 315 Bard, A. & Kock, M. 1994, A&A, 282, 1014 Barklem, P. S., Piskunov, N. & O'Mara, B. J. 2000, A&A, 363, 1091 Barnard A.J., Cooper J., & Smith E.W. 1974, JQSRT, 14, 1025 Bengston, R.D., & Miller, M.H. 1970, JOSA, 60, 1093 Berry, H.G., Bromander, J., Curtis, L.J. & Buchta, R. 1971, Phys. Scr., 3, 125 ´ Biemont E., Grevesse N., Faires L.M., et al. 1989, A&A, 209, 391 Bizzarri, A., Huber, M.C.E., Noels, A., et al. 1993, A&A, 273, 707 Blackwell, D. E., Booth, A. J., Petford, A. D. & Laming, J. M. 1989, MNRAS, 236, 235 Blackwell-Whitehead, R. J., Pickering, J. C., & Pearse, O. 2005, ApJS, 157, 402 Blanco, F., Botho, B. & Campos, J. 1995, Phys. Scr., 52, 628 Boesgaard, A. M., Heacox, W. D., Wolff, S. C., Borsenberger, J. & Praderie, F. 1982, ApJ, 259, 723 Breger, M. 1976, ApJS, 32, 7 Caliskan, H. & Adelman, S. J. 1997, MNRAS, 288, 501 Chen, L., Hou, J. L., & Wang, J. J. 2003, AJ, 125, 1397 Cunha, K., Hubeny, I., & Lanz, T. 2006, ApJ, 647, L143 Cutri, R. M., Skrutskie, M. F., van Dyk, S., et al., 2003, 2MASS All Sky Catalog of point sources. The IRSA 2MASS All-Sky Point Source Catalog, NASA/IPAC Infrared Science Archive. Donati, J.-F., Semel, M., Carter, B. D., Rees, D. E., & Collier Cameron, A. 1997, MNRAS, 291, 658 Dworetsky, M. M., & Budaj, J. 2000, MNRAS, 318, 1264 Fekel, F. C. 2003, PASP, 115, 807 Feldman, U., & Widing, K. G. 1990, ApJ, 363, 292 Fossati, L., Bagnulo, S., Monier, R., et al. 2007, A&A, 476, 911 Fossati, L., Bagnulo, S., Landstreet, J., et al. 2008, A&A, 483, 891 Fuhr, J.R., Martin, G.A. & Wiese, W.L. 1988. J. Phys. Chem. Ref. Data 17, Suppl. 4 Gigas, D. 1986, A&A, 165, 170 Gigas, D. 1988, A&A, 192, 264 Glagolevskij, Y. V. 1994, Bull. Special Astrophys. Obs., 38, 152 Gregg, M. D., Silva, D., Rayner, J. et al. 2004, Bulletin of the American Astronomical Society, The HST/STIS Next Generation Spectral Library, 36, 1496 Grevesse, N., Noels, A. & Sauval, A. J. 1996, Astronomical Society of the Pacific Conference Series, 99, 117 Hannaford, P., Lowe, R.M., Grevesse, N., Biemont, E. & Whaling, W. 1982, ApJ, 261, 736 Hauck, B., & Mermilliod, M. 1998, A&AS, 129, 431 Heacox, W. D. 1979, ApJS, 41, 675 Hibbert, A. 1988, Phys. Scr., 38, 37 Hill, G.M., & Landstreet, J.D. 1993, A&A, 276, 142 Hempel, M., & Holweger, H. 2003, A&A, 408, 1065 Holt, R. A., Scholl, T. J., & Rosner, S. D. 1999, MNRAS, 306, 107 Hubrig, S. & Castelli, F. 2001, A&A, 375, 963 ¨ Hubrig, S., Yudin, R. V., Scholler, M. & Pogodin, M. A. 2006, A&A, 446, 1089 Jamar, C., Macau-Hercot, D., Monfils, A. et al. 1976, Ultraviolet bright-star spectrophotometric catalogue. A compilation of absolute spectrophotometric data obtained with the Sky Survey Telescope (S2/68) on the European Astronomical Satellite TD-1 Johansson, S. 1978, Phys. Scr., 18, 217 Johnson, H. L. & Morgan, W. W. 1953, ApJ, 117, 313 Kelleher, D.E., & Podobedova, L.I. 2008, J. Phys. Chem. Ref. Data, 37, No.3, 1285 Kling, R. & Griesmann, U. 2000, ApJ, 531, 1173 Kling, R., Schnabel, R. & Griesmann, U. 2001, ApJS, 134, 173 Kocer, D., Adelman, S. J., Caliskan, H., Gulliver, A. F. & Gokmen Tektunali, H. 2003, A&A, 406, 975 Kochukhov, O. 2007, Spectrum synthesis for magnetic, chemically stratified stellar atmospheres, Physics of Magnetic Stars, 109, 118 Kochukhov, O., Tsymbal, V., Ryabchikova, T., Makaganyk, V., & Bagnulo S. 2006, A&A, 460, 831 Kontizas, E. & Theodossiou, E. 1980, MNRAS, 192, 745


18

Fossati et al.: The abundance analysis for normal A- and B-type stars.
¨ Shulyak, D., Tsymbal, V., Ryabchikova, T., Stutz, Ch., & Weiss, W. W. 2004, A&A, 428, 993 Sigut, T.A.A. 1999, ApJ, 519, 303. Singh, J. & Castelli, F. 1992, A&A, 253, 431 Smith, G. & Gallagher, A. 1966, Phys. Rev., 145, 26 Smith, K. C. & Dworetsky, M. M. 1993, A&A, 274, 335 Smith, G., O'Neil, J.A. 1975, A&A, 38, 1 Smith, G., & Raggett, D.St.J. 1981, J. Phys. B, 14, 4015 Smith, M.W. & Wiese, W.L. 1971, ApJS, 23, 103 Smith, G. 1988, J. Phys. B., 21, 2827 Smith, K. C. 1993, A&A, 276, 393 Smith, K. C. 1994, A&A, 291, 521 Sobeck, J. S., Lawler, J. E. & Sneden, C. 2007, ApJ, 667, 1267 Takeda, Y. 1994, PASJ, 46, 181 Takeda, Y. 2008, MNRAS, 388, 913 Theodosiou, E. 1989, Phys. Rev. A, 39, 4880 Theodossiou, E. & Danezis, E. 1991, A&AS, 183, 91 Tsymbal, V., Lyashko, D. & Weiss, W. W. 2003, IAU Symposium, 210, Ed. by Piskunov, N., Weiss, W. W. & Gray, D. F., 49 Wade, G. A., Ryabchikova, T. A., Bagnulo, S., & Piskunov, N. 2001, in Magnetic Fields Across the Hertzsprung-Russell Diagram, eds. G. Mathys, S. K. Solanki, and D. T. Wickramasinghe, ASP Conf. Proc., 248, 373 Wade, G. A., Donati, J.-F., Landstreet, J. D., & Shorlin, S. L. S. 2000, MNRAS, 313, 823 Wahlgren, G. M. & Hubrig, S. 2000, A&A, 362, L13 Warner, B. 1968a, MNRAS, 139, 115 Warner, B. 1968b, MNRAS, 140, 53 Wedemeyer, S. 2001, A&A, 373, 998 Wickliffe M.E. & Lawler J.E. 1997a, ApJS, 110, 163 Wiese, W.L., Smith, M.W., & Glennon, B.M. 1966, NSRDS-WSG 4 Wiese, W.L., Smith, M.W., & Miles, B.M. 1969, NSRDS-WSG 22 Wiese, W. L. & Fuhr, J. R. 2007, J. Phys. Chem. Ref. Data 36, 1287 ¨ ¨ Wilke, R. 2003, PhD thesis, Heinrich-Heine-Universitat, Dusseldorf Yoon, J., Peterson, D. M., Zagarello, R. J., Armstrong, J. T. & Pauls, T. 2008, ApJ, 681, 570

Kupka, F., Piskunov, N. E., Ryabchikova, T. A., Stempels, H. C., & Weiss, W. W. 1999, A&AS, 138, 119 Kurucz, R.L. & Peytremann, E. 1975, SAO Special Report 362 Kurucz, R.L. 1988, Trans. IAU, XXB, M. McNally, ed., Dordrecht: Kluwer, 168172 Kurucz, R. L. 1993, SYNTHE spectrum synthesis programs and line data Lacy, C. H. S., Fekel, F. C., Mathieu, R. D., et al. 1997, AJ, 113, 1088 Landstreet, J. D. 1998, A&A, 338, 1041 Lanz, T., Dimitrijevic, M.C. & Artru, M.-C. 1988, A&A, 192,249 Lanz, T., Cunha, K., Holtzman J., & Hubeny, I. 2008, ApJ, 678, 1342 Lawler, J.E. & Dakin, J.T. 1989, JOSA B6, 1457 Lemke, M. 1990, A&A, 225, 125 Leone, F. & Lanzafame, A.C. 1998, A&A, 330, 306 Lilly, R.A. 1976, JOSA, 66, 971 Ljung G., Nilsson H., Asplund M., & Johansson S. 2006, A&A456, 1181 ´ Ma´ iz-Apellaniz, J. 2005, PASP, 117, 615 ´ Ma´ iz-Apellaniz, J. 2006, AJ, 131, 1184 Malagnini, M. L., Faraggiana, R. & Morossi, C. 1983, A&A, 128, 375 Malfait, K., Bogaert, E. & Waelkens, C. 1998, A&A, 331, 211 Martin, G.A., Fuhr, J.R., & Wiese, W.L. 1988, J. Phys. Chem. Ref. Data, 17, Suppl.3 Mashonkina, L., Ryabchikova, T., & Ryabtsev, A. 2005, A&A, 441, 309 Matheron, P., Escarguel, A., Redon, R., Lesage, A. & Richou, J. 2001, JQSRT, 69, 535 Megessier, C. 1988, A&AS, 72, 551 Miles, B.M. & Wiese, W.L. 1969, WSG Technical Note 474 Miller, M.H., Roig, R.A. & Bengston, R.D. 1971. Phys. Rev. A, 4, 1709 Miller, M.H., Wilkerson, T.D., Roig, R.A. & Bengston, R.D. 1974, Phys. Rev. A, 9, 2312 Moon, T. T. & Dworetsky, M. M. 1985, MNRAS, 217, 305 Morel, T., & Butler, K. 2008, A&A, 487, 307 Morossi, C. & Malagnini, M. L. 1985, A&AS, 60, 365 Napiwotzki, R., Schoenberner, D. & Wenske, V. 1993, A&A, 268, 653 Netopil, M., Paunzen, E., Maitzen, H. M., North, P. & Hubrig, S. 2008, A&A, 491, 545 Nicolet, B. 1978, A&AS, 34, 1 Nieva, M. F. & Przybilla, N. 2008, A&A, 481, 199 Nilsson, H., Ljung, G., Lundberg, H., & Nielsen, K. E. 2006, A&A, 445, 1165 North, P. & Nicolet, B. 1990, A&A, 228, 78 O'Brian, T.R., & Lawler, J.E. 1991, Phys. Rev. A, 44, 7134 O'Brian, T.R., Wickliffe, M.E., Lawler, J.E., Whaling, W. & Brault, J.W. 1991, JOSA B8, 1185 Pamyatnykh, A. A. 1999, Acta Astronomica, 49, 119 Pickering, J. C., Thorne, A. P., & Perez, R. 2001 ApJS, 132, 403 Piskunov, N. E., Kupka, F., Ryabchikova, T. A., Weiss, W. W., & Jeffery, C. S. 1995, A&AS, 112, 525 Przybilla, N., Butler, K., Becker, S.R., Kudritzki, R.P., & Venn, K.A. 2000, A&A, 359, 1085 Przybilla, N., & Butler, K. 2001, A&A, 379, 955 Przybilla, N., Butler, K., & Kudritzki, R.P. 2001a, A&A, 379, 936 Przybilla, N., Butler, K., Becker, S.R., & Kudritzki, R.P. 2001b, A&A, 369, 1009 Przybilla, N., Nieva, M.-F., Butler, K. 2008, ApJ, 688, L103 Raassen, A.J.J., Pickering, J.C. & Uylings, P.H.M. 1998, A&AS, 130, 541 Raassen, A.J.J., & Uylings, P.H.M. 1998, A&A, 340, 300 Ralchenko, Yu., Kramida, A.E., Reader, J., & NIST ASD Team. 2008, NIST Atomic Spectra Database (version 3.1.5), available: http://physics.nist.gov/asd3, NIST, Gaithersburg, MD Rentzsch-Holm, I. 1996, A&A, 312, 966 Rentzsch-Holm, I. 1997, A&A, 317, 178 Roby, S. W. & Lambert, D. L. 1990, ApJS, 73, 67 Rufener, F. 1988, Catalogue of stars measured in the Geneva Observatory photometric system : 4 : 1988, Sauverny: Observatoire de Geneve Ryabchikova T.A., Hill, G.M., Landstreet J.D., Piskunov, N., & Sigut, T. A. A. 1994, MNRAS, 267, 697 Ryabchikova, T. A., Piskunov, N. E., Stempels, H. C., Kupka, F., & Weiss, W. W. 1999, Phis. Scr., T83, 162 Ryabchikova, T. A., Wade, G. A., & LeBlanc. F. 2003, in Modelling of Stellar Atmospheres, Proc. 210th IAU Symp., eds. N. Piskunov, W.W. Weiss, and D.F. Gray, p.301 Ryabchikova, T., Leone, F., & Kochukhov, O. 2005, A&A, 438, 973 Ryabchikova, T., Ryabtsev, A., Kochukhov, O., & Bagnulo, S. 2006, A&A, 456, 329 Sadakane, K. 1981, PASP, 93, 587 Sadakane, K., Jugaku, J. & Takada-Hidai, M. 1988, ApJ, 325, 776 Schiller, F. & Przybilla, N. 2008, A&A, 479, 849 Schulz-Gulde, E. 1969, JQSRT, 9, 13 Schlegel, D. J., Finkbeiner, D. P. & Davis, M. 1998, ApJ, 500, 525 Seaton, M.J., Mihalas, D., & Pradhan, A.K. 1994, MNRAS, 266, 805


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 1

Online Material


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 2

1.4

1.2

normalised flux

1

0.8

0.6

0.4

0.2 4330 4340 4350 4850 4860 4870
o

6550 6560 6570 6580

wavelength (A)
Fig. 8. Same as Fig. 2, but also for H and H.

1.4

1.2

normalised flux

1

0.8

0.6

0.4

0.2 4330 4340 4350 4850 4860 4870
o

6550 6560 6570 6580

wavelength (A)
Fig. 9. Same as Fig. 8, but for HD 145788.


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 3

1.4

1.2

normalised flux

1

0.8

0.6

0.4

0.2 4330 4340 4350 4850 4860 4870
o

6550 6560 6570 6580

wavelength (A)
Fig. 10. Same as Fig. 8, but for Cet.

Fig. 11. Comparison between LLMODELS theoretical fluxes calculated with the fundamental parameters and abundances derived for HD 145788, taking into account a reddening of E(B-V)=0.20 (full black line) and without taking into account reddening (dashed black line), with IUE calibrated fluxes (full blue line), Johnson UBV photometry (green squares), Geneva photometry (red circles) and 2MASS photmetry (violet triangles). The model fluxes were convolved to have approximately the same spectral resolution of the IUE fluxes (R 900).


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 4

1

0.95

normalised flux

0.9

0.85

observation Teff = 10400 Teff = 10600 Teff = 10200

4469

4470

4471

4472
o

4473

4474

wavelength (A)

Fig. 12. Comparison between the observed spectrum of the He I line at 4471 å for 21 Peg (black thick line) and three synthetic profiles calculated with three different T eff for the model atmosphere: 10200 K (green line), the adopted 10400 K (red line) and 10600 K (blue line).


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 5 Table 6. Comparison of the derived abundances with previous determinations for 21 Peg. Ion He I CI C II NI N II OI Ne I Na I Mg I Mg II Al I Al II Al III Si I Si II Si III P II S II Ca I Ca II Sc II Ti II V II Cr I Cr II Mn I Mn II Fe I Fe II Fe III Co II Ni I Ni II Zn I Sr II Y II Zr II Ba II Nd III R08 log(N /Ntot ) -1.11±0.04 -3.66±0.14 -3.65±0.05 -3.95±0.12 -3.90: -3.28±0.11 -3.76±0.05 -3.89 (non-LTE) -5.60: -4.42±0.12 -4.56±0.03 -5.89 -5.70±0.10 -4.95: -4.49±0.13 -4.26: -6.37±0.06 -4.86±0.13 -5.84±0.11 -5.98±0.08 -9.37±0.10 -7.23±0.09 -7.98±0.06 -6.29±0.09 -6.20±0.10 -6.54±0.21 -6.51±0.17 -4.52±0.13 -4.50±0.12 -4.60±0.06 -6.75±0.18 -5.71±0.05 -5.61±0.09 -6.96: -9.10: -9.76±0.15 -9.48±0.28 -9.19±0.06 -10.09±0.07 S81 log(N /Ntot ) -3.88 S934 log(N /Ntot ) D00 log(N /Ntot )

n 7 9 4 4 1 18 7 1 1 5 7 2 4 1 22 2 3 26 3 5 7 59 9 5 68 5 19 108 406 3 3 10 23 1 2 4 4 3 3

n 2

n

-3.60±0.16 -3.82 (non-LTE) -4.72±0.13 -6.05 7 2 -4.34±0.05 -5.84 -5.89±0.21 -4.85 2 -4.63±0.07 4 1 2 2

3 1

-5.76 -6.24 -9.47 -7.15±0.16 -7.78±0.16 -6.39±0.25 -6.08 -4.80±0.21 -4.79±0.18

1 1 3 26 6 17 2 32 23 -6.30±0.01 -6.39±0.10 -4.59±0.05 -6.01±0.01 5 3 10 3 4 3

-6.13 -9.18 -9.27 -9.09

5 2 1 1

-5.70±0.02 -7.29±0.11

In column 2, the meaning of a colon is the same as in Table 4. R08: this work; S81: Sadakane (1981); S934: Smith & Dworetsky (1993); Smith (1993, 1994); D00: Dworetsky & Budaj (2000).


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 6 Table 7. Comparison of the derived abundances with previous dterminations for Cet. Ion He I C II NI N II OI O II Ne I Na I Mg I Mg II Al I Al II Al III Si I Si II Si III P II P III S II Cl II Ar I Ar II Ca II Sc II Ti II Cr II Mn II Fe I Fe II Fe III Co II Ni II Sr II R08 log(N /Ntot ) n -0.97±0.04 6 -3.58±0.07 7 -4.03±0.13 10 -3.74±0.07 9 -3.06±0.14 9 -3.04: 2 -3.66±0.09 20 -3.86 (NLTE) 1 -5.23±0.07 3 -4.27: 2 -4.47±0.16 10 -5.57: 2 -5.73±0.27 8 -5.30±0.02 3 -4.80: 1 -4.41±0.20 31 -4.16: 2 -6.38±0.19 9 -6.19: 1 -4.78±0.16 31 -6.95: 2 -4.86: 2 -5.24±0.19 6 -5.77: 2 -9.31: 1 -7.42±0.08 11 -6.41±0.10 21 -6.50±0.09 3 -4.53±0.22 7 -4.58±0.14 186 -4.52±0.10 4 -6.93: 1 -5.76±0.19 17 -9.15: 2 mic = 1 km s-1 A91 log(N /Ntot ) -1.11±0.06 -3.81±0.07 -3.92±0.15 -3.34 S934 log(N /Ntot ) D00 log(N /Ntot ) A04 log(N /Ntot )

n 6 4 3 2

n

n

n

-3.29±0.03 -3.71±0.19 -3.86 (NLTE) 6 1 -4.57±0.02

4

-4.88 -4.56±0.08 -5.85 -5.36 -4.56±0.12 -4.99 -4.86±0.18

1 7 2 1 5 1 18

-4.28±0.05 -6.04 -5.79±0.10 -4.44±0.05

4 1 2 2

4

-4.44±0.28

4

-4.90

1

-5.72 -7.17±0.24 -6.58±0.19 -4.66±0.20 -4.82 -6.02

1 14 15 59 1 3 -6.00±0.2 -6.44±0.05 -4.55±0.05 -6.30±0.1 -5.80±0.2 mic = 0 km s- 5 3 10 3 4
1

-7.23±0.13 -6.58±0.08 -4.78±0.19 -5.89±0.02

7 10 29 2
1

mic = 0 km s-1

mic = 3 km s-

In column 2, the meaning of a colon is the same as in Table 4. R08: this work; A91: Adelman (1991); S934: Smith & Dworetsky (1993); Smith (1993, 1994); D00: Dworetsky & Budaj (2000); A04: Acke & Waelkens (2004)


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 7 Table 8. A collection of the experimental and theoretical transition probabilities and Stark widths for the observed Si II lines. Errors are given in parenthesis. Wavelength å 3853.665 3856.018 3862.595 4072.709 4075.452 4076.780 4128.054 4130.872 4130.894 4190.707 4198.134 4200.658 4621.418 4621.722 5041.023 5055.983 5056.316 5466.460 5469.451 5540.807 5576.661 5632.966 5639.477 5669.563 5688.817 5701.370 5800.454 5806.731 5827.752 5846.121 5867.480 5868.443 5957.559 5978.930 6347.109 6371.371 6660.532 6665.030 6671.841 6699.431 7848.816 7849.618 7849.722 excit eV 6.858 6.859 6.858 9.837 9.839 9.837 9.837 9.839 9.839 13.492 13.487 12.525 12.525 12.526 10.067 10.074 10.074 12.525 12.880 14.489 14.504 14.186 14.528 14.200 14.186 14.175 14.504 14.489 14.489 14.504 14.504 14.528 10.067 10.074 8.121 8.121 14.504 14.495 14.528 14.495 12.525 12.526 12.526 logg f BBC -1.39(06) -0.43(05) -0.80(06) log Wilke -5.12 -5.15 -5.14 -4.89 -4.89 -4.89 -4.92 -4.93 -5.25 -5.26 St LDA -5.15 -5.15 -5.15 -4.76 -4.76 -4.76 -4.88 -4.88 -4.88 -5.26 -5.26 -3.43 -3.69 -3.69 -4.78 -4.78 -4.78 -3.99

NIST -1.341 -0.406 -0.757 -2.700 -1.403 -1.700 0.359 -0.783 0.552 -0.889 -0.608 -0.453 0.029 0.523 -0.492 -0.237 -0.762 -0.818 0.286 0.126 -0.057

Barach

SG -1.32(06) -0.34(06) -0.83(06) -2.37(05) -1.40(05) -1.67(05) 0.45(04) 0.50(04)

BBCB -1.52 -0.56 -0.82

0.32 -0.82 0.48

Math -1.28(05) -0.36(06) -0.80(07) -2.70(10) -1.40(10) -1.70(10) 0.36(12) 0.55(12) -0.17(10) -1.46(11)

Wilke -1.53(11) -0.65(10) -0.92(10) -2.40(16) -1.64(16) -1.75(16) 0.22(08) 0.38(08) -0.33(11) -0.60(11)

AJPP -1.44 -0.49 -0.75 -1.89 -0.95 -1.20 0.38 -0.77 0.53 -0.35 -0.61

0.29(03) 0.59(03) -0.36(03) -0.83(08) -0.47(08) -0.62(08) -0.07(06) 0.32(06) 0.08(06) -0.10(06) -0.04(08) -0.10(08) -0.91(18) -0.51(08) 0.42(06) -0.30(02) 0.00(02) 0.18(05) -0.06(05)

0.09(04) 0.46(04) -0.49(04)

0.03(07) 0.52(08)

0.18(03) 0.48(03)

0.19 0.44 -0.51 -0.80 -0.49 -0.69 -0.32 0.25 0.08 -0.10 -0.17 -0.20 -0.90 -0.70 -0.20 0.20 -0.33 -0.03 0.17 -0.13 0.24 -0.16 0.52 -0.16 0.34 -0.80 0.50

-4.84 -4.79 -4.79

-1.20(10) -0.47(10) -0.82(06) -0.30(07) 0.28(07) 0.13(06) -0.06(07) -0.16(07) -0.13(07) -0.79(07) -0.29(07) -0.22(03) 0.08(03) 0.20(04) -0.03(03) 0.41(07) -0.25(08) -0.07(12)

-0.12(12) -0.11(12) -1.00(10) -0.57(12) 0.40(14)

-0.62(13) -0.82(09) -0.18(09) 0.12(08) 0.00(09) -0.28(09) -0.16(11) -0.18(10) -1.16(11) -0.33(14) 0.50(11) -0.29(06) 0.03(06) 0.29(08) -0.02(08) 0.23(09) -0.18(09) 0.46(09)

-5.46 -5.44 -5.44 -5.53 -5.50 -5.53 -5.37 -5.47 -5.45 -5.35 -5.36 -5.02 -5.01 -5.31 -5.32 -5.51 -5.54 -5.58 -4.84 -4.84 -5.05 -5.05

-0.225 0.084 0.149 -0.082 0.162 -0.240 0.409 -0.247 0.316 -0.831 0.470

0.30 -0.00

0.12(07) -0.04(07)

-4.34 -4.34 -4.34

NIST ­ Kelleher & Podobedova (2008) ­ critical compilation; Barach ­ Barach (1970) ­ pulsed arc; SG ­ Schilz-Gulde (1969) ­ arc emission spectra; BBCB ­ Berry et al. (1971) ­ beam-foil lifetime measurements and theortical branching ratios; BBC ­ Blanco et al. (1995) ­ laser-induced plasma; Math ­ Matheron et al. (2001) ­ laser-induced plasma; Wilke ­ Wilke (2003) ­ laser-induced plasma, Stark widths, shifts and transition probabilities; AJPP ­ Artru et al. (1981) ­ theoretical calculations; LDA ­ Lanz et al. (1988) ­ Stark widths, semi-classical calculations.


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 8 Table 9. Linelist of the lines used for the abundance analysis.

Element Wavelength å He I 4026.1850 4026.1870 4026.1870 4026.1980 4026.3580 4026.1990 4026.3580 4120.8110 4120.8240 4120.9910 4168.9670 4387.9290 4437.5510 4471.4690 4471.4730 4471.4730 4471.4850 4471.4880 4471.6820 4713.1390 4713.1560 4713.3760 4921.9310 CI 4029.4130 4762.5282 4770.0263 4771.7458 4775.8947 4932.0490 5052.1670 5380.3370 6014.8342 7111.4694 7113.1790 7116.9879 7119.6572 C II 3918.9702 4266.9991 4267.2590 5145.1666 6578.0530 6582.8817 7236.4158 7237.1656 NI 6482.6986 6484.8081 6644.9650 6722.6141 7423.6453 7442.3003 7468.3132 8216.3313 8680.2858

excit eV 20.964 20.964 20.964 20.964 20.964 20.964 20.964 20.964 20.964 20.964 21.218 21.218 21.218 20.964 20.964 20.964 20.964 20.964 20.964 20.964 20.964 20.964 21.218 7.488 7.483 7.483 7.488 7.488 7.685 7.685 7.685 8.643 8.640 8.647 8.647 8.643 16.332 18.046 18.046 20.710 14.449 14.449 16.333 16.333 11.764 11.758 11.764 11.845 10.326 10.330 10.336 10.336 10.336

logg f

HD 145788 EQW abundance må dex S S S S S S S -1.13

EQW må S S S S S S S S S S S S S S S S S S S S S S 1.92 1.77 4.98 2.38 2.82 4.31 1.48 2.10 6.06 5.03 4.59 4.60 11.10 12.74 S S

21 Peg abundance dex -1.15

EQW må

Cet abundance dex

Ref logg f

-2.620(005) -1.450(005) -0.700(005) -1.450(005) -1.320(005) -0.970(005) -1.320(005) -1.743(041) -1.963(041) -2.433(041) -2.338(005) -0.883(005) -2.034(041) -2.198(005) -1.028(005) -0.278(005) -1.028(005) -0.548(005) -0.898(005) -1.233(041) -1.453(041) -1.923(041) -0.435(005) -2.190 -2.335(100) -2.437(100) -1.866(100) -2.304(100) -1.658(041) -1.303(041) -1.616(041) -1.585(180) -1.086(041) -0.774(041) -0.907(041) -1.149(041) -0.533(041) 0.562(068) 0.717(068) 0.189(068) -0.026(041) -0.328(041) 0.298(041) -0.656(041) -0.510(100) -0.674(100) -0.859(140) -0.714(100) -0.706(025) -0.384(025) -0.189(025) 0.133(025) 0.346(025)

-1.05

S S S S S S S S S S S S S S

-0.99 -1.01 -0.98 -0.95 -0.90

S S S S S S S S S S 5.4 19.0 12.4 22.7 13.3

-1.15

-1.10 -1.08 -1.08

-1.05 -1.05 -3.63 -3.36 -3.65 -3.71 -3.64

-1.13 -1.10 -3.61 -3.55 -3.65 -3.55 -4.00 -4.16 -4.30 -3.61 -3.55 -3.94 -3.85 -3.61 -3.61 -3.67 -3.70 -3.60

-0.99

WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG WSG KP NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08

5.5 8.3

-3.68 -3.58

S S S 2.5 S S 11.7 3.1

-3.60 -3.63 -3.60 -3.60 -3.58 -3.64 -3.45

NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08

3.77 4.89 2.70 2.61

-4.10 -3.82 -3.89 -3.99 7.9 11.8 17.2 39.0 45.9 -3.85 -3.97 -3.95 -4.11 -4.17


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 9 Table 9. continued.

Element Wavelength å 8683.4050 8686.1542 8703.2496 8711.7073 8718.8354 N II 3994.9970 4447.0300 4607.1494 4613.8676 4630.5432 4643.0899 5005.1538 5666.6274 5679.5540 OI 3947.2928 4772.9155 4773.7553 4802.1377 4802.9775 4967.3739 4968.7935 5019.2901 5020.2198 5330.7372 5435.1793 5435.7791 5436.8588 5512.7686 5958.3893 5958.5793 6106.2650 6155.9864 6156.7761 6158.1858 6453.6065 6454.4462 O II 4641.8102 Ne I 5656.6590 5764.4190 5820.1560 5852.4880 5881.8950 5975.5340 6029.9968 6074.3370 6096.1631 6143.0630 6163.5940 6217.2810 6266.4950 6334.4280 6382.9914 6402.2461 6506.5278

excit eV 10.330 10.326 10.326 10.330 10.336 18.497 20.409 18.462 18.466 18.483 18.483 20.666 18.466 18.483 9.146 10.741 10.741 10.741 10.741 10.740 10.741 10.741 10.741 10.741 10.740 10.741 10.741 10.989 10.989 10.989 14.047 10.740 10.741 10.741 10.740 10.741 22.979 18.613 18.555 18.576 16.848 16.619 16.619 16.671 16.671 16.671 16.619 16.715 16.619 16.715 16.619 16.671 16.619 16.671

logg f 0.086(025) -0.305(025) -0.322(025) -0.234(025) -0.335(025) 0.276(100) 0.285(068) -0.507(025) -0.665(025) 0.094(025) -0.359(025) 0.592(041) -0.045(005) 0.250(005) -2.096(041) -1.742(068) -1.567(100) -2.122(068) -1.976(068) -1.630(068) -1.275(068) -1.871(068) -1.725(068) -0.870(068) -1.766(068) -1.544(068) -1.398(068) -1.942(068) -1.742(068) -1.471(068) -0.638(100) -0.664(025) -0.443(025) -0.296(025) -1.288(068) -1.066(068) 0.054(041) -0.83(13) -0.370 -0.581 -0.455(041) -0.747(041) -1.249(041) -1.037(041) -0.477(068) -0.297(068) -0.098(041) -0.603(041) -0.955(041) -0.357(041) -0.315(041) -0.230(041) 0.345(041) -0.021(041)

HD 145788 EQW abundance må dex

EQW må

21 Peg abundance dex

EQW må 26.1 25.3 18.1 23.7 19.0 S 4.2 2.6 1.7 6.4 3.6 3.5 1.9 3.9

Cet abundance dex -4.28 -3.91 -4.07 -4.02 -4.03 -3.85 -3.66 -3.73 -3.78 -3.79 -3.67 -3.65 -3.83 -3.72

Ref logg f NIST08 NIST08 NIST08 NIST08 NIST08 WSG WSG WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF WF BM L L NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08

S

-3.90

14.9

-3.29

8.7 14.7

-2.99 -3.11

S 4.64 5.26 2.45 3.11 5.10 11.10 2.54 4.52 27.90 4.18 7.73 9.99

-3.45 -3.32 -3.44 -3.22 -3.26 -3.37 -3.37 -3.43 -3.32 -3.28 -3.27 -3.21 -3.24

2.6 6.1 3.7 S S

-3.23 -3.06 -2.78 -3.11 -2.88 -3.12 -3.13 -3.08 -3.13 -3.08 -3.55 -3.72 -3.79 -3.71 -3.57 -3.76 -3.62 -3.71 -3.63 -3.62 -3.65 -3.59 -3.56 -3.57 -3.66 -3.51

46.3 65.3 13.8 20.0

-3.11 -3.08 -3.12 -3.15

1.70 34.49 48.82 10.46 17.15

-3.02 -3.22 -3.20 -3.20 -3.17

25.5 37.4 45.2 8.0 11.2 2.2 1.9 3.5 1.8 10.2 9.2 2.3 4.4 10.1 14.9 20.2 8.4 5.1 13.8 15.0 14.2 34.4

2.36 1.40

-3.77 -3.83

3.40 4.33 2.08 7.26 4.14

-3.70 -3.76 -3.82 -3.77 -3.69


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 10 Table 9. continued.

Element Wavelength å 6598.9528 7032.4130 7488.8712 7535.7739 Na I 5688.2051 5895.9240 8183.2550 8194.8240 Mg I 4167.2710 4702.9912 5167.3208 5172.6841 5183.6040 5528.4048 Mg II 3848.2110 3848.3410 3850.3860 4384.6372 4390.5718 4427.9941 4433.9878 6787.8550 7877.0540 7896.0420 7896.3660 8213.9870 8234.6360 Al I 3944.0060 3961.5200 Al II 4663.0459 5593.2998 6226.1948 6231.6210 6231.7500 6243.2030 6243.3672 7042.0830 7063.6820 7471.4100 Al III 4529.1890 5696.6040 5722.7300 Si I 3905.5230 Si II 3853.6650 3856.0180 3862.5950 4072.7090 4075.4520 4076.7800 4128.0540

excit eV 16.848 16.619 18.382 18.382 2.104 0.000 2.102 2.104 4.346 4.346 2.709 2.712 2.717 4.346 8.864 8.864 8.864 9.996 9.999 9.996 9.999 12.085 9.996 9.999 9.999 9.996 9.999 0.000 0.014 10.598 13.257 13.071 13.073 13.073 13.077 13.077 11.317 11.317 13.649 17.818 15.642 15.642 1.909 6.858 6.859 6.858 9.837 9.839 9.837 9.837

logg f -0.342(041) -0.249(041) 0.167(100) 0.04(18) -0.450(005) -0.184(005) 0.230(005) 0.490(005) -0.745(068) -0.440(025) -0.870(025) -0.393(025) -0.167(005) -0.498(025) -1.580(005) -2.534(025) -1.842(005) -0.776(005) -0.523(005) -1.208(005) -0.907(005) -0.911(025) 0.391(005) -0.308(025) 0.643(005) -0.271(005) 0.032(005) -0.623(068) -0.323(068) -0.284(010) 0.41(18) 0.05(18) -0.08(18) 0.40(18) -0.08(18) 0.67(18) 0.35(07) -0.35(07) 0.78(18) 0.671(068) 0.235(041) -0.068(041) -1.041(023) -1.341(100) -0.406(068) -0.757(068) -2.367(100) -1.403(068) -1.700(068) 0.359(041)

HD 145788 EQW abundance må dex

EQW må

21 Peg abundance dex

EQW må 9.9 12.9 3.6 4.4 24.6 7.6 11.4

Cet abundance dex -3.65 -3.56 -3.86 -3.62 -5.30 -5.17 -5.21 -4.24 -4.29 -4.51 -4.41 -4.48 -4.49 -4.40 -4.26 -4.40 -4.30 -4.76 -4.73 -5.57 -5.57 -5.65 -5.86 -5.73 -5.80 -5.80 -5.28 -5.49 -6.19 -5.28 -5.30 -5.32 -4.80: -4.50 -4.35 -4.30 -4.56 -4.48 -4.58

Ref logg f NIST08 NIST08 NIST08 NIST08 WSM WSM WSM WSM NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 SW SW WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM NIST08 NIST08 NIST08 BL KP08 KP08 KP08 KP08 KP08 KP08 KP08

3.21

-5.60

36.2 48.8 123.9 140.0 46.6 S S S 29.8 42.2

-4.32 -4.40 -3.96 -3.87 -4.32 -4.24 -4.24 -4.31 -4.37

16.67 40.42 69.62 85.63 14.39 S S S 39.51 60.83 20.81 33.15

-4.53 -4.46 -4.36 -4.24 -4.52 -4.58 -4.51 -4.55 -4.60 -4.53 -4.56 -4.56

2.4 18.4 S S S 19.0 30.4 3.5 86.1 S S 48.9 68.1 S S 35.1 10.0 6.1 S S S S 47.0 15.4 4.3 5.8 3.9 2.2

S S 32.6 11.9

-5.70 -5.70 -5.22 -5.01

S S 25.28 4.80 2.77

-5.89 -5.89 -5.60 -5.83 -5.65

12.8

-5.12

8.93

-5.72

S S S S S S

-4.75: -4.30 -4.20 -4.16 -4.30 -4.50

S S S S 3.44 S S

-4.95 -4.55 -4.40 -4.35 -4.45 -4.52 -4.52

S S S S 22.1 15.3 S


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 11 Table 9. continued.

Element Wavelength å 4130.8720 4130.8940 4190.7070 4198.1340 4200.6582 4621.4180 4621.7222 5041.0239 5055.9839 5056.3169 5466.4609 5469.4512 5540.8070 5576.6610 5632.9660 5639.4771 5669.5630 5688.8170 5701.3700 5800.4540 5806.7310 5827.7520 5846.1210 5867.4800 5868.4438 5957.5591 5978.9302 6371.3711 6660.5320 6665.0300 6671.8410 6699.4310 7848.8160 7849.6180 7849.7220 Si III 4552.6220 4567.8400 4574.7570 P II 4420.7120 4602.0690 5296.0770 5344.7290 5409.7220 5425.8800 6034.0390 6043.0840 6459.9450 P III 4222.1980 S II 3923.4450 3998.7590 4028.7500 4142.2590 4153.0680 4162.6650

excit eV 9.839 9.839 13.492 13.487 12.525 12.525 12.526 10.067 10.074 10.074 12.525 12.880 14.489 14.504 14.186 14.528 14.200 14.186 14.175 14.504 14.489 14.489 14.504 14.504 14.528 10.067 10.074 8.121 14.504 14.495 14.528 14.495 12.525 12.526 12.526 19.016 19.016 19.016 11.021 12.853 10.802 10.737 10.755 10.802 10.737 10.802 10.935 14.610 16.198 16.246 15.944 15.848 15.899 15.944

logg f -0.783(140) 0.552(041) -0.350 -0.607 -0.889(200) -0.608(180) -0.453(180) 0.029(041) 0.523(041) -0.492(041) -0.237(100) -0.762(180) -0.796 -0.493 -0.818(068) -0.318 0.286(025) 0.126(041) -0.057(041) -0.170 -0.203 -0.903 -0.699 -0.203 0.199 -0.225(025) 0.084(025) -0.082(068) 0.162(068) -0.240(068) 0.409(025) -0.247(041) 0.316(041) -0.831(100) 0.470(041) 0.292(025) 0.068(025) -0.409(041) -0.478 0.799(080) -0.134 -0.329 -0.368 0.241 -0.209 0.384 0.161 0.21(18) 0.44(20) 0.06(20) 0.00(20) 0.24(20) 0.62(20) 0.78(20)

HD 145788 EQW abundance må dex S -4.45 S

EQW må S S 3.51 7.02 9.20 73.29 86.78 41.20 10.73 2.68

21 Peg abundance dex -4.52

EQW må S S 11.8 8.5

Cet abundance dex -4.60 -4.49 -4.35

Ref logg f KP08 KP08 AJPP81 AJPP81 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 AJPP81 AJPP81 KP08 AJPP81 KP08 KP08 KP08 AJPP81 AJPP81 AJPP81 AJPP81 AJPP81 AJPP81 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 KP08 H MRB H H H H H H H NIST08 WSM WSM WSM WSM WSM WSM

85.8 98.2 41.4

-4.06 -4.39 -4.27

-4.73 -4.56 -4.59 -4.28 -4.58 -4.40 -4.43 -4.43

S S S 6.2 1.2 3.7 2.5 10.6 7.8 6.4 4.7

-4.29 -4.49 -4.39 -4.43 -4.23 -4.19 -4.54 -4.56 -4.48 -4.40 -4.01 -3.96 -4.20 -4.31 -4.42 -4.45 -4.20 -4.30 -4.12 -4.20 -4.89 -4.90

1.99

-4.20

2.63

-4.14 2.1 3.7 6.7 12.9 60.5 74.3 S S 14.8 S 12.5 S S 9.3 4.7 2.1 3.4 3.2 1.6 1.2 5.6 3.7 8.5 1.7 1.8

62.8 138.4

-4.23 -4.08

3.87 35.37 48.33 110.50 4.18

-4.30 -4.57 -4.56 -4.31 -4.20

2.70 1.20

-4.13 -4.39

-4.16 -4.17 -6.26 -6.56 -6.31 -6.46 -6.53 -6.41 -6.06 -6.21 -6.63 -6.19 -4.83 -4.80 -4.74 -4.77

1.82 1.21 2.43

-6.44 -6.33 -6.35

1.44 3.22 3.39 8.2 -4.45 7.36

-5.09 -4.71 -4.88 -4.89

S 7.3 18.9 21.4


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 12 Table 9. continued.

Element Wavelength å 4230.9430 4269.7250 4294.4020 4463.5810 4486.6340 4716.2710 4792.0070 4815.5520 4917.1980 4925.3430 4942.4730 5009.5670 5014.0420 5032.4340 5103.3320 5201.0270 5212.6200 5320.7230 5345.7120 5428.6550 5432.7970 5453.8550 5509.7050 5526.2430 5556.0230 5564.9580 5578.8700 5606.1510 5616.6330 5639.9770 5640.3460 5647.0200 5664.7730 5819.2540 6312.6850 Cl II 4794.5560 4810.0700 Ar I 8103.6920 8115.3110 Ar II 4277.5280 4426.0010 4430.1890 4609.5670 4735.9060 4806.0210 Ca I 4226.7280 4302.5278 4434.9570 4454.7788 5588.7490 6439.0752 Ca II 3933.6630 5001.4790

excit eV 17.446 16.092 16.135 15.944 15.867 13.617 16.135 13.672 14.003 13.584 13.584 13.617 14.068 13.672 13.672 15.068 15.068 15.068 15.068 13.584 13.617 13.672 13.617 13.701 13.617 13.672 13.677 13.733 13.660 14.068 13.701 14.003 13.660 14.068 14.234 13.375 13.375 11.624 11.548 18.454 16.749 16.813 18.454 16.644 16.644 0.000 1.899 1.886 1.899 2.526 2.526 0.000 7.505

logg f 0.56 -0.12(20) 0.58(20) -0.02(18) -0.40(18) -0.41(14) -0.12(20) 0.09(14) -0.32(18) -0.46(18) -0.96(18) -0.28(18) 0.10(18) 0.27(18) -0.11(18) 0.09(18) 0.32(18) 0.49(18) 0.35(18) -0.13(18) 0.26(18) 0.48(18) -0.14(18) -0.53(18) -0.99(18) -0.32(18) -0.51(18) 0.31(18) -0.64(18) 0.28(18) 0.15(20) 0.04(18) -0.25(18) -0.76(14) -0.14(18) 0.455(100) 0.292(100) -0.130(010) 0.360(010) -0.060(041) 0.158(005) -0.174(041) 0.304(005) -0.108(005) 0.210(005) 0.244 0.292(010) -0.007(020) 0.258(020) 0.358(011) 0.390(010) 0.105(027) -0.507

HD 145788 EQW abundance må dex

EQW må 4.71 2.21 3.34 6.21 2.34 3.86 2.80 4.87 6.32 2.01 1.86 2.46 2.89 3.57 5.51 7.27 3.15

21 Peg abundance dex -4.82 -4.63 -4.85 -4.91 -4.87 -4.62 -4.94 -4.79 -4.97 -5.21 -4.81 -4.89 -4.76 -4.78 -4.90 -4.91 -4.79

EQW må 4.0 5.5 7.2 3.1 11.30 2.6 18.7 9.6 3.9 12.6 16.9 7.1 13.3 13.3 28.7 11.5 3.6 2.6 9.0 4.4 16.7 4.2 13.0 6.9 3.0 3.6 0.5 0.8 5.6 8.2 1.0 2.2 1.9 1.0 1.6 1.3

Cet abundance dex -5.11 -4.72 -4.64 -4.75 -4.76 -4.90 -4.83 -4.74 -4.77 -4.72 -4.69 -5.20 -4.76 -4.71 -4.63 -4.76 -5.01 -4.74 -4.70 -4.92 -4.86 -4.81 -4.65 -4.91 -4.64 -4.96 -7.15 -6.76 -4.69 -5.03 -5.04 -5.37 -5.09 -5.23 -5.15 -5.54

Ref logg f Mult WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM WSM MWRB WSM WSM WSM WSM WSM WSM NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 NIST08 SG S SN SN SR SR

10.9

-4.27

4.31 4.50 3.60 3.51 1.64

-4.97 -4.74 -4.91 -4.68 -4.94

62.8 11.4 17.7 12.9 14.2 S 23.4

-5.63 -5.44 -5.48 -5.38 -5.35 -5.70 -5.57

10.44 1.40 1.77

-5.95 -5.85 -5.72

S 5.56

-5.85 -5.99

S 3.6

-5.85 -5.69

T TB


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 13 Table 9. continued.

Element Wavelength å 5019.9712 5021.1382 5285.2661 5307.2241 6456.8750 Sc II 4246.8218 4294.7671 4374.4570 4400.3892 4415.5571 5031.0210 5526.7900 Ti II 4012.3835 4028.3384 4053.8208 4161.5293 4163.6436 4171.9038 4174.0718 4287.8735 4290.2148 4300.0420 4301.9224 4374.8159 4386.8467 4391.0249 4394.0586 4395.0312 4395.8389 4399.7651 4409.5195 4411.0723 4411.9292 4417.7134 4418.3311 4421.9385 4441.7290 4443.8008 4450.4819 4464.4487 4468.5068 4469.1499 4488.3247 4501.2700 4518.3315 4529.4800 4533.9600 4544.0161 4549.6216 4563.7573 4568.3237 4571.9712 4657.2002 4708.6621 4763.8833 4764.5239

excit eV 7.515 7.515 7.505 7.515 8.438 0.315 0.605 0.618 0.605 0.595 1.357 1.768 0.574 1.892 1.893 1.084 2.590 2.598 2.598 1.080 1.165 1.180 1.161 2.061 2.598 1.231 1.221 1.084 1.243 1.237 1.231 3.095 1.224 1.165 1.237 2.061 1.180 1.080 1.084 1.161 1.131 1.084 3.124 1.116 1.080 1.572 1.237 1.243 1.584 1.221 1.224 1.572 1.243 1.237 1.221 1.237

logg f -0.247 -1.207 -1.147 -0.848 0.405 0.242(020) -1.391(040) -0.418(074) -0.536(043) -0.668(045) -0.400(047) 0.024(050) -1.84(04) -0.96(03) -1.13(03) -2.09(04) -0.13(03) -0.29(03) -1.26(04) -1.79(02) -0.85(03) -0.44(03) -1.15(03) -1.61(03) -0.96(04) -2.28(05) -1.78(04) -0.54(02) -1.93(04) -1.19(03) -1.84(08) -0.67(03) -2.52(08) -1.19(02) -1.97(04) -1.66(04) -2.33(04) -0.72(02) -1.52(02) -1.81(02) -0.60(02) -2.33(05) -0.51(03) -0.77(02) -2.56 -1.64(03) -0.53(02) -2.58(07) -0.11(03) -0.69(02) -2.94 -0.32(03) -2.24(03) -2.34(05) -2.36(06) -2.95

HD 145788 EQW abundance må dex 36.0 -5.54 8.6 -5.38 10.5 -5.34 30.0 69.1 25.7 19.2 13.6 21.2 -5.71 -8.92 -8.87 -8.91 -8.92 -8.87

EQW må 10.14 1.15 2.21 19.40 0.89 4.84 3.95 2.83 1.89 3.31 17.99 19.95 15.14 5.46 34.44 4.54 8.99 38.52 55.58 24.27 4.74 10.18 3.61 8.35 55.42 7.53 24.00 6.57 11.11 2.09 24.81 5.49 4.77 3.75 7.52 16.61 9.60 51.62 3.30 13.46 45.64 2.38 51.37 1.71 64.47 40.98 1.02 56.42 3.59 2.87 2.50 1.25

21 Peg abundance dex -5.96 -6.05 -6.05 -9.50 -9.17 -9.38 -9.37 -9.39 -9.42 -9.36 -7.16 -7.24 -7.23 -7.24 -7.30 -7.30 -7.31 -7.30 -7.27 -7.35 -7.24 -7.22 -7.17 -7.28 -7.23 -7.16 -7.28 -7.33 -7.19 -7.18 -7.30 -7.28 -7.19 -7.13 -7.26 -7.26 -7.22 -7.25 -7.24 -7.24 -7.24 -7.17 -7.27 -7.20 -7.11 -7.38 -7.08 -7.15 -7.21 -7.22 -7.27 -6.98

EQW må

Cet abundance dex

Ref logg f TB TB TB TB TB

2.5

-9.31

LD LD LD LD LD LD LD PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP BHN PTP PTP PTP Kur00 PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP

49.9 22.0 73.7 65.6 10.7 32.0 82.9 103.5 31.4 16.2 31.6 103.8 26.1 61.5 28.7 7.7 62.7 24.0 18.3 13.2 92.7 51.6 32.4 96.5 35.5 92.7 10.8 29.8 97.6 9.6 87.7 106.2 16.7 13.0

-6.72 -6.85 -6.79 -6.82 -7.17 -6.94 -6.70 -6.59 -6.87 -6.76 -6.87 -6.52 -6.83 -6.85 -6.93 -6.90 -6.87 -6.85 -6.82 -6.85 -6.64 -6.82 -6.87 -6.64 -6.93 -6.59 -6.44 -6.85 -6.63 -6.72 -6.73 -6.44 -6.69 -6.82

5.8 4.5 4.5 10.4 2.2

-7.42 -7.37 -7.48 -7.48 -7.51

11.5

-7.38

7.0

-7.44

5.6 10.6

-7.52 -7.41


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 14 Table 9. continued.

Element Wavelength å 4779.9849 4798.5317 4805.0850 4911.1948 5010.2104 5013.6860 5072.2866 5129.1562 5154.0679 5185.9019 5188.6870 5211.5303 5226.5386 5262.1411 5268.6074 5336.7861 5381.0210 5418.7676 V II 4002.9360 4005.7050 4023.3780 4035.6220 4036.7770 4202.3550 4225.2130 4528.4850 4564.5790 4600.1660 Cr I 4274.7969 4289.7168 5204.5112 5206.0371 5208.4248 Cr II 4003.2830 4012.4961 4017.9629 4037.9719 4049.0969 4051.9299 4054.0759 4056.0559 4072.5610 4132.4189 4145.7808 4172.5908 4179.4209 4195.4170 4207.3628 4215.7378 4236.3789 4242.3638 4252.6318 4254.5220 4256.1079 4261.9131

excit eV 2.048 1.080 2.061 3.124 3.095 1.582 3.124 1.892 1.566 1.893 1.582 2.590 1.566 1.582 2.598 1.582 1.566 1.582 1.428 1.817 1.805 1.793 1.476 1.704 2.026 2.276 2.268 2.265 0.000 0.000 0.941 0.941 0.941 6.484 5.662 5.330 6.487 6.484 3.104 3.105 5.662 3.714 3.758 5.319 3.105 3.827 5.319 3.827 3.104 3.104 3.871 3.858 5.871 6.487 3.864

logg f -1.26(05) -2.68(08) -0.96(06) -0.61(03) -1.29(06) -2.19(06) -1.06(04) -1.24(03) -1.75(03) -1.49(03) -1.05(02) -1.16(05) -1.26(02) -2.25(07) -1.67 -1.59(03) -1.92(05) -2.00(05) -1.447(041) -0.522(027) -0.689(027) -0.767(230) -1.594(020) -1.523(117) -1.463(032) -1.098(029) -1.393(067) -1.523(119) -0.231(041) -0.361(041) -0.208(041) 0.019(041) 0.158(041) -0.723 -1.085 -2.279 -0.679 -1.032 -2.360 -2.478 -1.812 -2.632 -2.568 -1.106 -2.943 -1.996 -1.753 -2.731 -3.168 -3.606 -1.363 -2.054 -1.294 -1.520 -1.560

HD 145788 EQW abundance må dex 34.8 -6.84 50.1 33.2 12.2 17.7 36.3 28.3 32.7 58.7 14.2 48.0 -6.83 -6.88 -6.81 -6.81 -6.93 -6.79 -6.75 -6.87 -7.15 -6.88

EQW må 10.19 2.44 17.19 10.83 3.37 2.72 4.44 11.03 6.56 10.18 20.67 3.42 2.64 2.34 8.66 4.20 3.45 3.99 15.73 11.23 3.10 2.97 1.80 2.72 1.88 1.52 5.40 3.29 2.30 3.64 5.27 9.41 17.26 1.10 11.23 6.19 12.29 8.12 1.87 5.28 5.23 15.42 3.15 16.14 3.92 3.69 1.44

21 Peg abundance dex -7.24 -7.05 -7.25 -7.25 -7.15 -7.20 -7.23 -7.31 -7.25 -7.10 -7.33 -7.56 -7.16 -7.21 -7.26 -7.28 -7.28 -7.99 -8.02 -8.04 -7.94 -7.91 -8.02 -8.06 -7.94 -7.91 -6.33 -6.43 -6.23 -6.25 -6.22 -6.33 -6.07 -6.37 -6.29 -6.23 -6.35 -6.44 -6.43 -6.17 -6.22 -6.29 -6.43 -6.18 -6.15 -6.13 -6.12 -6.22 -6.23 -6.22 -6.16

EQW må 1.3

Cet abundance dex -7.22

Ref logg f RHL PTP RHL PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP PTP K00 PTP PTP PTP BGF BGF BGF BGF BGF BGF BGF BGF BGF BGF MFW MFW MFW MFW MFW RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

2.2

-7.41

15.2 15.5

-6.97 -6.88

33.5 33.8 7.7

-7.64 -7.56 -7.74

12.1 24.5 17.2 8.7 18.5

-7.27 -6.15 -6.23 -6.19 -6.03

19.5 27.4

-6.03 -6.08

3.3 2.8

-6.49 -6.44

28.0 33.5

-6.02 -5.86

1.5

-6.30

64.8 27.0 51.9

-5.78 -5.93 -5.88

13.03 6.99 2.16 30.94

6.7

-6.48


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 15 Table 9. continued.

Element Wavelength å 4269.2769 4275.5669 4284.1880 4306.9160 4465.7310 4511.7749 4539.5952 4554.9878 4558.6499 4565.7402 4588.1992 4592.0488 4616.6289 4618.8032 4634.0698 4697.5981 4723.3398 4736.9990 4779.0532 4812.3369 4824.1270 4901.6230 4912.4619 4920.2300 4952.7939 5153.4990 5232.4961 5237.3291 5243.4570 5246.7681 5274.9639 5279.8760 5280.0542 5305.8530 5308.4082 5313.5630 5334.8691 5369.3560 5407.6040 5420.9219 5478.3652 5502.0669 5508.6060 5620.6309 5678.3901 6053.4658 6089.6318 6129.2261 Mn I 4030.7529 4034.4829 4041.3550 4783.4268 4823.5239 Mn II 4200.2700 4205.3770 4206.3672

excit eV 3.854 3.858 3.854 5.873 6.487 6.484 4.042 4.071 4.073 4.042 4.071 4.074 4.072 4.074 4.072 5.670 4.168 4.156 5.670 3.864 3.871 6.487 6.484 3.887 6.282 3.758 4.071 4.073 4.042 3.714 4.071 4.073 4.074 3.827 4.071 4.074 4.072 3.871 3.827 3.758 4.178 4.168 4.156 6.487 6.484 4.745 6.487 4.750 0.000 0.000 2.114 2.298 2.319 6.185 1.809 5.397

logg f -2.201 -1.736 -1.897 -1.602 -1.279 -1.375 -2.394 -1.491 -0.662 -1.982 -0.845 -1.473 -1.576 -1.084 -1.236 -1.913 -2.784 -2.872 -2.409 -2.125 -1.085 -1.141 -1.262 -3.161 -1.573 -2.500 -2.360 -1.350 -2.764 -2.466 -1.559 -2.112 -2.316 -2.160 -2.058 -1.779 -1.826 -3.045 -2.459 -2.458 -1.968 -2.117 -2.252 -1.395 -1.496 -2.219 -1.445 -2.511 -0.470(068) -0.811(068) 0.285(068) 0.042(041) 0.144(041) -1.741 -3.440(113) -1.553(041)

HD 145788 EQW abundance må dex 22.7 -5.89 42.2 -5.91 31.1 -6.00

15.1 49.9 97.8 27.0 87.5 53.0 47.2 72.1 66.3

-5.82 -5.88 -5.54 -5.90 -5.62 -5.82 -5.85 -5.77 -5.76

24.4 9.7 13.9 13.0 57.2 12.8 46.1

-5.92 -5.78 -5.92 -5.91 -5.85 -5.93 -5.87

14.4 17.3 25.5 23.1 15.4

-5.90 -5.74 -5.85 -5.77 -5.87

EQW må 9.91 23.31 17.13 3.93 2.59 2.37 5.57 28.81 61.28 12.95 54.80 29.19 25.53 44.97 38.59 2.75 1.25 1.39 2.04 11.25 47.81 5.71 4.32 1.19 2.85 5.66 6.30 33.00 2.28 5.85 26.66 8.64 6.85 10.73 12.14 18.36 16.02 1.81 6.91 5.74 11.92 8.69 6.55 3.40 2.53 3.05 2.49 2.01 3.02 2.29 2.65 1.99 1.10 2.06 4.02 4.22

21 Peg abundance dex -6.22 -6.19 -6.23 -6.18 -6.38 -6.32 -6.21 -6.17 -6.11 -6.19 -6.12 -6.17 -6.17 -6.15 -6.17 -6.13 -6.42 -6.29 -5.77 -6.22 -6.18 -6.12 -6.13 -6.22 -6.12 -6.24 -6.15 -6.18 -6.23 -6.27 -6.14 -6.24 -6.15 -6.22 -6.12 -6.16 -6.19 -6.14 -6.14 -6.26 -6.15 -6.17 -6.18 -6.06 -6.10 -6.22 -6.12 -6.11 -6.81 -6.60 -6.44 -6.25 -6.60 -6.27 -6.62 -6.54

EQW må 4.4 3.7

Cet abundance dex -6.50 -6.43

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU MFW MFW MFW MFW MFW K88 KG KSG

0.8 7.3 27.9 3.3 24.4 6.7 5.5 15.1 11.6

-6.53 -6.41 -6.48 -6.31 -6.39 -6.47 -6.46 -6.44 -6.43

2.2

-6.43

6.8 7.8

-6.58 -6.30

4.1 1.1 4.4

-6.34 -6.41 -6.10

10.6 10.1

-6.50 -6.35

14.2

-6.02


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 16 Table 9. continued.

Element Wavelength å 4251.7168 4252.9629 4283.7690 4292.2368 4478.6372 4518.9561 4727.8408 4730.3950 4738.2900 4755.7271 4764.7280 4784.6250 4791.7822 4806.8232 5559.0469 5570.5391 6122.4502 Fe I 3998.0527 4005.2419 4006.3108 4007.2722 4017.1484 4021.8665 4030.4885 4045.8125 4062.4409 4063.2759 4063.5942 4067.9778 4071.7380 4132.0581 4132.8994 4134.6777 4143.4146 4143.8682 4153.8999 4154.4985 4154.8057 4156.7988 4157.7803 4158.7925 4172.1221 4175.6362 4181.7549 4187.0391 4191.4307 4196.2085 4199.0952 4202.0293 4217.5454 4219.3604 4222.2129 4224.1719 4225.4541 4233.6030 4235.9370 4247.4253

excit eV 6.185 6.186 5.373 5.380 6.645 6.645 5.370 5.373 5.380 5.397 5.398 6.573 6.185 5.418 6.185 6.177 10.184 2.692 1.557 3.267 2.758 3.047 2.758 3.211 1.485 2.845 3.368 1.557 3.211 1.608 1.608 2.845 2.831 3.047 1.557 3.396 2.831 3.368 2.831 3.417 3.430 3.251 2.845 2.831 2.449 2.469 3.396 3.047 1.485 3.430 3.573 2.449 3.368 3.417 2.482 2.425 3.368

logg f -1.058 -1.138 -2.204 -1.544(041) -0.950 -1.329 -2.017 -2.147 -1.876(061) -1.242 -1.351 -1.505 -1.715 -1.848(061) -1.318 -1.444 0.950 -0.910(041) -0.610(025) -0.99(10) -1.276(037) -1.063(025) -0.729(029) -0.555(037) 0.280(025) -0.862(045) -0.804 0.062(021) -0.472(037) -0.022(025) -0.675(021) -1.006(037) -0.649(045) -0.204(037) -0.511(021) -0.321(021) -0.688(041) -0.400(029) -0.809(037) -0.403(021) -0.67(10) -0.893(057) -0.827(037) -0.371(037) -0.548(025) -0.666(041) -0.696(025) 0.155(037) -0.708(025) -0.484(025) 0.000(021) -0.967(025) -0.506(025) -0.510(025) -0.604(025) -0.341(025) -0.239(037)

HD 145788 EQW abundance må dex 11.6 -6.19 10.8 -6.18

EQW må 4.11 4.15 2.52 3.22 2.80 1.81 2.19 2.02 3.11 5.76 1.51 1.01 1.75 2.72 1.94 1.49 2.91 14.97 1.66 1.95 1.72

21 Peg abundance dex -6.64 -6.55 -6.14 -6.69 -6.68 -6.49 -6.38 -6.29 -6.36 -6.58 -6.42 -6.59 -6.62 -6.50 -6.53 -6.87 -4.50 -4.62 -4.36 -4.28 -4.39 -4.65 -4.64 -4.47 -4.41 -4.67 -4.42 -4.68 -4.60 -4.56 -4.48 -4.50 -4.60 -4.42 -4.62 -4.39 -4.53 -4.50 -4.48 -4.60 -4.57 -4.63 -4.55 -4.50 -4.60 -4.55 -4.58 -4.59

EQW må 2.1 1.4

Cet abundance dex -6.52 -6.58

Ref logg f K88 K88 K88 KSG K88 K88 K88 K88 KSG K88 K88 K88 K88 KSG K88 K88 K88 BWL MFW MFW BWL BWL BWL BWL MFW BWL K88 BWL BWL MFW BWL BWL BWL BWL BWL BWL BWL BWL BWL BWL MFW BWL BWL BWL MFW BWL BWL BWL MFW BWL BWL MFW BWL BWL MFW MFW BWL

17.4

-6.20 3.1 -6.40

20.6 96.1 17.5 79.5

-4.16 -3.95 -4.07 -4.20 2.47 46.20 2.86 1.97 36.30 4.77 31.65 13.17 1.72 4.54 4.54 3.26 4.69 2.91 1.96 2.73 7.88 8.66 5.22 14.87 2.35 6.66 3.52 2.38 6.61 11.02 4.71 7.4 5.0 5.5 -4.73 -4.66 -4.51

31.3 52.2 18.4 14.6 18.5 15.5 17.8 10.5 15.6 30.7 32.2 20.3 14.2 40.3 47.6 11.3 23.2 17.6 14.7 27.1 38.2 19.9

-4.26 -4.37 -4.27 -4.36 -4.21 -4.21 -4.20 -4.20 -4.18 -4.23 -4.25 -4.41 -4.04 -4.42 -4.33 -4.35 -4.35 -4.21 -4.23 -4.29 -4.34 -4.33


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 17 Table 9. continued.

Element Wavelength å 4250.1196 4250.7871 4260.4746 4271.1538 4271.7607 4299.2349 4383.5449 4404.7505 4415.1226 4422.5679 4442.3389 4443.1943 4447.7173 4454.3809 4466.5518 4469.3755 4482.2529 4484.2197 4494.5630 4528.6143 4556.1260 4667.4531 4668.1318 4678.8457 4707.2744 4736.7734 4789.6509 4903.3101 4918.9941 4920.5029 4957.2988 4957.5967 4966.0889 4973.1021 4982.5239 4983.2671 4983.8652 4985.2529 5005.7109 5006.1191 5007.2930 5014.9424 5022.2354 5049.8198 5074.7485 5125.1172 5133.6885 5137.3823 5139.2515 5139.4629 5162.2729 5165.4072 5192.3442 5195.4678 5215.1807 5217.3892 5232.9404 5263.3062

excit eV 2.469 1.557 2.399 2.449 1.485 2.425 1.485 1.557 1.608 2.845 2.198 2.858 2.223 2.831 2.831 3.654 2.223 3.602 2.198 2.176 3.602 3.602 3.266 3.602 3.241 3.211 3.546 2.882 2.865 2.832 2.851 2.808 3.332 3.960 4.103 4.154 4.103 3.928 3.884 2.832 4.103 3.943 3.984 2.279 4.220 4.220 4.178 4.178 2.998 2.940 4.178 4.220 2.998 4.220 3.266 3.211 2.940 3.266

logg f -0.405(025) -0.714(021) 0.109(029) -0.349(025) -0.164(025) -0.405(021) 0.200(025) -0.142(025) -0.615(025) -1.115(025) -1.255(025) -1.043(025) -1.342(025) -1.299(025) -0.600(057) -0.477(068) -1.482(083) -0.864(037) -1.136(025) -0.822(025) -0.787(029) -0.751(041) -1.295 -0.833(041) -1.08(20) -0.752(041) -0.958(037) -0.926(025) -0.342(029) 0.068(025) -0.408(029) 0.233(021) -0.871(061) -0.95(07) 0.144 -0.158 -0.068 -0.560(037) -0.180 -0.638(025) -0.210 -0.303(053) -0.53(07) -1.355(029) -0.20(07) -0.14(18) 0.14(18) -0.40(07) -0.741(025) -0.509(025) 0.02(18) -0.035 -0.421(025) -0.002 -0.871(029) -1.070(029) -0.058(021) -0.879(029)

HD 145788 EQW abundance må dex 33.9 -4.34 47.0 -4.29 56.3 -4.42 35.8 -4.37 75.4 -4.16 33.8 -4.38 90.5 -4.10 77.1 -4.11 49.4 -4.32 15.1 14.2 23.4 10.1 18.8 12.8 9.7 10.6 9.0 -4.16 -4.09 -4.19 -3.93 -4.16 -3.93 -3.72 -3.94 -3.98

EQW må 9.30 12.90 21.64 11.23 32.47 8.76 45.27 31.24 15.68 1.77 3.31 2.33 0.76 4.99 3.44 3.35 1.65 3.35 6.37 2.36 1.46 1.36 2.21 1.27 2.36 6.71 16.48 6.39 21.42 3.19 1.95 2.90 1.83 3.19 4.31 1.55 2.31 1.70

21 Peg abundance dex -4.58 -4.60 -4.66 -4.55 -4.60 -4.64 -4.61 -4.61 -4.57 -4.45 -4.38 -4.39 -4.65 -4.50 -4.34 -4.13 -4.31 -4.50 -4.52 -4.23 -4.48 -4.43 -4.51 -4.37 -4.49 -4.59 -4.55 -4.56 -4.57 -4.75 -4.64 -4.58 -4.39 -4.54 -4.53 -4.72 -4.54 -4.42

EQW må 2.7 3.0

Cet abundance dex -4.61 -4.71

Ref logg f MFW BWL BWL,BK MFW MFW BWL MFW MFW MFW BWL MFW BWL MFW BWL BWL BWL BWL BWL MFW MFW BWL BWL K88 BWL MFW BWL BWL BWL BWL BWL BWL BWL BWL MFW K88 K88 K88 BWL K88 BWL,BK K88 BWL MFW BWL MFW MFW MFW MFW BWL BWL MFW K88 BWL K88 BK BKK BWL BKK

26.0 47.0 20.1 60.2 7.0 7.9 13.7 8.3 10.3 17.8 9.6 9.1 12.0 10.9 22.9 13.6 18.5 9.3 21.5 10.4 8.0

-4.37 -4.33 -4.47 -4.22 -4.27 -3.76 -4.49 -4.41 -4.43 -4.16 -4.32 -4.27 -4.15 -4.26 -4.16 -3.91 -4.17 -4.44 -4.33 -4.42 -4.07

5.55 2.26 3.16 6.76 3.92 5.37 5.50 11.83 2.61

-4.44 -4.31 -4.47 -4.38 -4.48 -4.55 -3.93 -4.55 -4.27


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 18 Table 9. continued.

Element Wavelength å 5266.5552 5269.5376 5281.7905 5283.6211 5324.1792 5328.0386 5339.9292 5364.8711 5367.4668 5369.9619 5371.4897 5383.3691 5393.1675 5400.5024 5410.9097 5415.1992 5424.0684 5429.6968 5434.5239 5445.0425 5446.9170 5462.9531 5463.2764 5572.8423 5586.7559 5602.9453 5615.6440 5862.3530 6230.7231 6400.0010 Fe II 3922.0040 3938.9700 4012.7439 4024.5471 4031.4419 4041.6411 4044.0120 4048.8320 4052.4751 4057.4610 4060.7419 4061.7820 4138.2080 4138.4048 4147.2720 4167.2990 4173.4610 4177.6919 4178.8618 4182.6890 4183.2002 4199.4912 4200.5210 4202.5220 4202.8589 4205.5952 4235.3940

excit eV 2.998 0.859 3.038 3.241 3.211 0.915 3.266 4.445 4.415 4.371 0.958 4.312 3.241 4.371 4.473 4.386 4.320 0.958 1.011 4.386 0.990 4.473 4.434 3.396 3.368 3.430 3.332 4.549 2.559 3.602 9.126 5.911 10.987 4.495 4.732 5.569 5.571 5.569 9.836 7.274 9.836 5.956 4.732 2.828 4.616 11.196 2.583 2.544 2.583 4.732 2.642 11.149 11.167 6.807 6.807 11.207 11.149

logg f -0.386(029) -1.321(025) -0.834(041) -0.432(029) -0.103(029) -1.466(025) -0.647(029) 0.228(049) 0.443(029) 0.536(021) -1.645(025) 0.645(021) -0.715(029) -0.16(18) 0.398(021) 0.642(021) 0.52(18) -1.879(025) -2.122(025) -0.02(18) -1.914(025) -0.156 0.11(07) -0.275(029) -0.120(033) -0.850(029) 0.050(029) -0.058 -1.281(025) -0.290(029) -1.106 -1.932 -0.562 -2.439 -3.162 -3.376 -2.671 -2.381 -1.276 -1.680 -1.416 -2.927 -3.494 -4.316 -3.788 -0.557 -2.617 -3.449 -2.535 -3.960 -5.090 -0.226 -0.303 -2.360 -2.665 -0.302 -0.758

HD 145788 EQW abundance må dex 22.5 -4.33 36.1 -4.36 11.2 -4.24 16.5 -4.32 18.6 16.0 20.0 24.4 22.8 29.9 9.8 14.1 20.3 31.6 33.6 11.6 14.4 17.3 19.7 24.9 27.9 9.1 15.0 -4.02 -4.29 -4.39 -4.39 -4.29 -4.40 -4.31 -4.01 -4.31 -4.32 -4.19 -4.15 -4.13 -4.16 -4.28 -4.31 -4.43 -4.16 -4.27

EQW må 5.62 8.37 2.16 3.52 7.14 6.56 3.35 3.26 5.23 5.96 4.14 10.10 2.24 2.69 4.40 8.21 8.51 3.15 1.57 2.29 2.14 2.22 2.34 4.23 1.52 7.11 1.81 2.10 2.93

21 Peg abundance dex -4.56 -4.63 -4.53 -4.60 -4.60 -4.57 -4.39 -4.62 -4.63 -4.69 -4.59 -4.57 -4.52 -4.36 -4.64 -4.63 -4.53 -4.48 -4.52 -4.57 -4.60 -4.39 -4.66 -4.57 -4.44 -4.69 -4.53 -4.34 -4.58

EQW må

Cet abundance dex

Ref logg f BWL MFW BWL BKK BKK MFW BKK BWL BWL,BK BWL MFW BWL BKK MFW BWL BWL MFW MFW MFW MFW BWL K88 MFW BKK BWL,BKK BK BKK K88 MFW BKK RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

3.5

-4.12

1.4

-4.39

S S 1.62 30.01 6.91 2.65 16.05 1.78 14.0 -3.75 1.41 5.17 4.74 5.16 2.29 29.40 57.71 0.96 1.53 2.10 3.56 2.45 5.08 2.72 -4.80 -4.47 -4.51 -4.32 -4.41 -4.56 -4.52 -4.27 -4.36 -4.48 -4.45 -4.51 -4.59 -4.61 -4.36 -4.82 -4.58 -4.44 -4.38 -4.24 6.2 11.4 8.0 3.5

-4.54 -4.54

-4.39 -4.38 -4.50 -4.23

88.3 9.8

-4.14 -4.03

3.9 38.4 14.5 42.8 6.8

-4.28 -4.70 -4.56 -4.67 -4.37


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 19 Table 9. continued.

Element Wavelength å 4250.4370 4258.1538 4258.3398 4263.8691 4273.3262 4278.1592 4286.2798 4296.5718 4303.1758 4369.4110 4384.0942 4384.3188 4385.3872 4402.8770 4413.6011 4416.8301 4418.9570 4431.6050 4444.2988 4445.2612 4446.2368 4448.5210 4449.6162 4450.3052 4451.5508 4451.9751 4453.2051 4455.2661 4460.8979 4461.4390 4461.7061 4466.6660 4472.9292 4487.4971 4489.1831 4491.4048 4493.5290 4499.2842 4499.6880 4504.3428 4507.1021 4508.2881 4512.0562 4512.3281 4515.3389 4515.6089 4520.2241 4522.6338 4526.4038 4526.5942 4534.1680 4541.5239 4549.1919 4549.4741 4555.8931 4556.3921 4559.5542 4563.1270

excit eV 7.684 2.704 2.642 7.693 2.704 2.692 7.708 2.704 2.704 2.778 6.226 2.657 2.778 6.138 2.676 2.778 7.946 7.940 6.219 11.207 5.956 11.149 7.929 6.223 6.138 11.255 7.684 6.226 11.291 2.583 6.226 6.217 2.844 7.693 2.828 2.856 7.920 7.790 7.693 6.219 7.773 2.856 7.684 7.708 2.844 6.226 2.807 2.844 7.806 5.569 2.856 2.856 5.911 2.828 2.828 7.693 7.804 7.806

logg f -1.719 -3.478 -4.301 -1.691 -3.303 -3.951 -1.715 -2.933 -2.511 -3.584 -2.579 -3.684 -2.582 -2.563 -4.185 -2.602 -1.848 -1.785 -3.274 -0.677 -2.776 -0.595 -1.699 -2.984 -1.907 -0.742 -2.020 -2.000 -0.810 -4.366 -2.065 -3.125 -3.531 -2.138 -2.971 -2.756 -1.562 -2.461 -1.676 -3.250 -1.760 -2.349 -2.203 -2.370 -2.540 -2.719 -2.617 -2.169 -2.177 -3.582 -3.364 -2.973 -1.767 -2.016 -2.421 -2.044 -2.447 -2.389

HD 145788 EQW abundance må dex 46.3 15.1 50.4 15.6 65.9 -4.18 -4.01 -4.26 -3.96 -4.28

EQW må 6.56 21.22 5.69 7.27 27.48 13.34 8.65 40.34 55.97 5.17 19.76 52.33 7.95 7.18 52.09 3.81 4.39 1.78 1.58 6.25 2.19 5.52 2.09 18.31 1.80 2.16 13.29 1.43 6.74 12.04 0.69 19.70 3.60 40.10 45.69 1.07 4.52 1.23 4.89 60.28 3.28 1.52 53.64 3.99 51.83 66.30 2.08 1.80 26.24 37.52 27.36 73.86 2.59 1.11 1.13

21 Peg abundance dex -4.49 -4.63 -4.55 -4.46 -4.63 -4.44 -4.34 -4.65 -4.61 -4.45 -4.50 -4.61 -4.30 -4.54 -4.60 -4.47 -4.47 -4.25 -4.51 -4.30 -4.47 -4.45 -4.47 -4.48 -4.36 -4.69 -4.54 -4.38 -4.43 -4.53 -4.82 -4.55 -4.33 -4.55 -4.58 -4.51 -4.68 -4.44 -4.52 -4.57 -4.30 -4.48 -4.58 -4.42 -4.57 -4.57 -4.48 -4.26 -4.51 -4.60 -4.44 -4.51 -4.57 -4.49 -4.54

EQW må 5.6 6.0 13.7 5.9 24.9 39.9 7.0

Cet abundance dex -4.45 -4.42 -4.67 -4.42 -4.68 -4.71 -4.70

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

79.7 14.8 80.0 6.2

-4.14 -3.96 -4.21 -4.17

36.7 2.2 2.8 3.5 3.8 13.8 8.4 9.1

-4.67 -4.63 -4.60 -4.37 -4.54 -4.48 -4.61 -4.50

9.1

-4.13

20.0 13.9 21.4 33.6 65.5 72.7 7.1 94.7 85.7 82.5 101.9

-4.29 -4.17 -4.18 -4.32 -4.18 -4.21 -4.40 -4.07 -4.09 -4.12 -4.08

23.9 29.7 4.4 4.1 2.7 44.1 38.7 35.6 51.0

-4.61 -4.65 -4.60 -4.62 -4.70 -4.69 -4.65 -4.66 -4.70

63.9 93.2

-4.21 -4.04

22.2 42.9

-4.64 -4.66


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 20 Table 9. continued.

Element Wavelength å 4576.3398 4579.5269 4580.0630 4582.8350 4583.8369 4591.0039 4595.6821 4596.0151 4598.4941 4601.3779 4605.3701 4610.5942 4619.6279 4620.5210 4625.4810 4625.8931 4628.7861 4629.3390 4631.8730 4635.3159 4638.0498 4640.8120 4648.9438 4652.2158 4656.9810 4663.7080 4666.7578 4670.1821 4720.1489 4731.4531 4804.7178 4820.8340 4824.8379 4825.7358 4826.6831 4908.1509 4913.2920 4923.9268 4924.9209 4937.0781 4942.1768 4948.0962 4948.7930 4951.5840 4952.6572 4953.9868 4958.8218 4969.3652 4976.0059 4977.0352 4977.9229 4984.4878 4990.5088 4991.1260 4991.4399 4993.3579 4999.1802 5000.7432

excit eV 2.844 6.226 2.583 2.844 2.807 7.845 2.856 6.226 7.804 2.891 7.869 5.571 7.708 2.828 7.880 5.956 7.845 2.807 7.869 5.956 7.708 7.708 2.583 7.880 2.891 2.891 2.828 2.583 3.197 2.891 7.708 10.308 8.145 2.635 10.288 10.329 10.288 2.891 2.844 10.308 10.308 10.308 10.348 10.308 10.348 5.571 10.379 10.360 9.100 10.360 10.329 10.329 10.329 2.778 10.273 2.807 10.273 2.778

logg f -2.976 -2.343 -3.904 -3.224 -1.867 -2.264 -4.583 -1.956 -1.536 -4.519 -2.286 -3.673 -1.965 -3.315 -2.131 -2.549 -1.700 -2.478 -1.945 -1.578 -1.536 -1.737 -4.565 -2.269 -3.643 -3.889 -3.368 -4.073 -4.822 -3.127 -2.458 -0.723 -2.172 -5.052 -0.500 -0.272 0.050 -1.504 -4.962 -1.387 -1.234 -0.218 -0.031 0.211 -0.621 -2.815 -0.762 -0.830 -1.599 -0.039 -0.599 0.078 0.195 -4.541 -0.589 -3.684 -0.435 -4.578

HD 145788 EQW abundance må dex 63.9 -4.21 13.5 -4.16 52.5 122.8 26.1 13.4 -4.22 -3.99 -4.15 -4.13

45.9 16.0

-4.28 -4.00

EQW må 37.02 9.69 12.82 29.65 76.69 1.96 2.59 18.02 7.52 3.10 2.37 1.05 2.57 25.57 2.67 7.77 6.10 3.19 31.97 9.72 4.29 3.42 1.83 15.20 8.85 25.67 8.94 2.07 31.46 1.64 3.18 1.49 1.65 4.77 6.94 11.86 99.11 1.16 1.74 0.83 6.58 9.00 13.77 3.46 5.43 3.25 2.21 1.35 9.50 2.89 10.85 13.12 3.35 4.13 15.07 4.72 2.14

21 Peg abundance dex -4.61 -4.36 -4.56 -4.56 -4.59 -4.39 -4.51 -4.38 -4.50 -4.48 -4.27 -4.40 -4.63 -4.59 -4.36 -4.40 -4.42 -4.47 -4.46 -4.41 -4.62 -4.55 -4.40 -4.56 -4.60 -4.54 -4.58 -4.19 -4.57 -4.32 -4.48 -4.43 -4.36 -4.52 -4.52 -4.55 -4.32 -4.49 -4.08 -4.57 -4.60 -4.60 -4.60 -4.50 -4.48 -4.38 -4.50 -4.56 -4.55 -4.62 -4.60 -4.60 -4.47 -4.48 -4.56 -4.56 -4.63

EQW må 21.7 7.6 5.4 15.7 62.3 14.8 5.9

Cet abundance dex -4.66 -4.31 -4.59 -4.60 -4.69 -4.35 -4.54

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

13.3 5.8 3.7 41.6 26.6 6.9 3.0

-4.61 -4.36 -4.57 -4.65 -4.47 -4.50 -4.69

48.1 9.7 20.3 46.0 54.1

-4.15 -4.15 -4.26 -4.22 -4.25

4.3 12.9 17.1 2.9 5.0 6.2 12.0

-4.57 -4.57 -4.63 -4.57 -4.54 -4.64 -4.64

13.1 14.8 146.0

-4.04 -4.31 -3.93

11.9 15.6 16.2 9.6

-4.15 -4.15 -4.39 -4.20

6.0 7.2 14.8 2.5 3.0 3.7 8.9 3.1 11.3 13.9 9.7 4.9

-4.72 -4.80 -4.66 -4.71 -4.56 -4.37 -4.67 -4.64 -4.67 -4.67 -4.41 -4.61

18.9 16.7 29.0

-4.13 -4.34 -4.28


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 21 Table 9. continued.

Element Wavelength å 5001.9590 5004.1948 5004.7729 5006.8408 5007.4468 5007.7388 5009.0220 5010.0601 5011.0288 5018.4399 5019.4619 5021.5942 5022.4199 5022.7920 5026.8062 5030.6299 5031.8979 5032.7119 5035.7080 5036.7178 5036.9199 5045.1138 5047.6411 5060.2568 5061.7178 5062.9351 5067.8931 5070.8989 5073.5620 5075.7642 5081.9009 5083.5059 5086.3062 5089.2139 5089.4932 5094.9058 5097.2710 5098.6851 5107.1099 5113.0040 5117.0342 5119.3408 5120.3521 5127.8662 5132.6689 5136.8018 5140.6919 5141.3701 5143.8799 5144.3550 5146.1270 5148.9072 5149.4648 5150.4888 5150.9409 5154.4092 5157.2822 5159.9121

excit eV 10.273 10.273 10.381 10.379 10.381 10.288 10.348 10.381 10.381 2.891 5.569 10.288 10.348 10.288 10.308 10.288 10.414 10.391 10.288 10.391 2.828 10.308 10.308 10.448 10.308 10.308 10.329 10.308 6.807 10.455 10.379 10.348 10.414 10.329 10.455 10.467 10.379 10.448 10.348 10.391 10.431 10.391 2.828 5.571 2.807 2.844 10.414 10.467 10.448 10.467 2.828 10.414 10.448 10.448 2.856 2.844 10.455 10.414

logg f 0.916 0.504 -1.128 -0.362 -0.460 -0.282 -0.530 -0.694 -1.183 -1.345 -2.784 -0.191 -0.073 -0.092 -0.444 0.431 -0.833 0.077 0.632 -0.565 -4.732 -0.002 -0.235 -0.650 0.284 -1.113 -0.078 0.268 -2.889 0.184 -1.062 -0.752 -0.419 0.013 -0.407 -0.720 0.315 -0.489 -0.826 -0.527 -0.039 -0.669 -4.256 -2.451 -4.094 -4.356 -1.190 -0.773 -0.205 0.307 -4.079 -0.417 0.554 -0.078 -4.483 -4.269 -0.173 -0.920

HD 145788 EQW abundance må dex 28.5 -4.28

EQW må

21 Peg abundance dex -4.57 -4.34 -4.45 -4.45 -4.48 -4.49 -4.71 -4.46 -4.35 -4.46 -4.57 -4.61 -4.48 -4.35 -4.62 -4.37 -4.58 -4.59 -4.51 -4.44 -4.73 -4.30 -4.66 -4.18 -4.71 -4.60 -4.30 -4.56 -4.22 -4.64 -4.57 -4.67 -5.05 -4.52 -4.43 -4.40 -4.68 -4.52 -4.51 -4.48 -4.50 -4.52 -4.05 -4.41 -4.51 -4.68 -4.59 -4.62 -4.52 -4.52 -4.57

EQW må 37.6

Cet abundance dex -4.66

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

10.0

-4.19

153.4

-3.99

13.6 20.9 13.3 31.5 15.6 18.7 24.1 20.7 23.1 15.5

-4.19 -4.43 -4.32 -4.31 -4.19 -3.84 -4.17 -3.90 -4.17 -4.29

21.61 1.55 6.18 5.06 7.38 4.16 1.80 1.06 105.97 6.05 7.31 7.89 10.11 6.72 18.16 2.67 10.45 24.21 3.56 2.19 7.55 4.31 13.73 2.42 6.66 14.46 1.69 12.19 2.29 1.86 4.04 8.33 1.45 2.28 4.49 2.60

3.3 88.3

-4.66 -4.45

19.4 26.0 9.6 3.6 15.0 7.0 17.6 12.7

-4.71 -4.67 -4.69 -4.45 -4.72 -4.76 -4.59 -4.65

17.7

-4.10

17.2 4.3 4.5 5.3 2.5 2.0 4.6

-4.63 -4.53 -4.49 -4.88 -4.62 -4.57 -4.36

13.0 15.1 19.3 18.0 11.4 9.5 17.6 12.9 22.5 9.3 27.7 14.9

-4.18 -4.08 -4.17 -4.27 -4.11 -4.18 -4.30 -3.81 -4.39 -4.32 -3.70 -3.94

6.70 2.69 5.21 10.89 7.58 4.09 2.34 2.51 6.60 11.78 6.36 6.70 3.04 6.82 1.35

5.7 9.8

-4.67 -4.88

6.9

-4.70


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 22 Table 9. continued.

Element Wavelength å 5160.8389 5161.1841 5166.5552 5169.0332 5170.7769 5175.3921 5177.0200 5178.3711 5180.3140 5186.8730 5194.8921 5197.5771 5199.1221 5200.8042 5203.6382 5213.5400 5214.4888 5215.3491 5218.8418 5219.6992 5219.9258 5222.3608 5223.7998 5224.4111 5228.8960 5231.9072 5232.7871 5234.6250 5236.6128 5237.9502 5243.1919 5245.4551 5247.9521 5253.6411 5254.3999 5254.9292 5256.9380 5257.1221 5257.8940 5260.2588 5262.3169 5262.9561 5263.4829 5264.8120 5265.9810 5268.4092 5270.0269 5271.1000 5272.3970 5276.0020 5278.9380 5284.1089 5291.6660 5298.8599 5303.3950 5306.1802 5311.8999 5313.1060

excit eV 5.569 2.856 10.455 2.891 10.455 7.727 10.379 10.431 10.391 10.467 10.467 3.230 10.379 10.391 10.391 10.523 10.503 10.379 10.381 10.480 10.523 10.519 10.379 10.414 10.448 10.531 10.381 3.221 10.455 10.448 8.259 10.455 10.531 10.452 10.500 3.230 2.891 10.500 10.455 10.419 10.545 10.500 10.467 3.230 10.419 10.500 10.503 10.503 5.956 3.199 5.911 2.891 10.480 10.519 8.185 10.523 10.545 10.519

logg f -2.559 -4.573 -0.045 -1.250 -0.330 -2.354 -0.197 -0.334 -0.088 -0.194 -0.108 -2.348 0.121 -0.036 -0.115 -0.761 -0.593 0.000 -0.165 -1.039 -0.551 -0.281 -0.503 -0.428 -0.300 -0.651 -0.082 -2.279 -0.676 0.104 -1.698 -0.543 0.550 -0.134 -0.461 -3.336 -4.182 0.156 -0.526 1.069 -0.368 -0.874 -0.955 -3.133 -0.871 -0.964 -0.197 -1.160 -2.009 -2.213 -2.677 -3.195 0.544 -0.405 -1.534 0.044 -1.023 -0.720

HD 145788 EQW abundance må dex 15.4 -4.19 5.9 -4.21 10.9 -4.25 160.8 -4.02 10.6 -3.97 9.4 13.9 7.0 92.2 15.6 11.1 -4.22 -4.09 -4.34 -3.90 -4.24 -4.20

EQW må 9.67

21 Peg abundance dex -4.43 -4.58 -4.35 -4.50 -4.63 -4.58 -4.71 -4.51 -4.75 -4.53 -4.34 -4.60 -4.88 -4.50 -4.62 -4.47 -4.57 -4.53 -4.60 -4.52 -4.57 -4.61 -4.71 -4.53 -4.43 -4.54 -4.39 -4.49 -4.61 -4.46 -4.46 -4.68 -4.54 -4.68 -4.44 -4.52 -4.67 -4.54 -4.54 -4.55 -4.37 -4.32 -4.45 -4.27 -4.72 -4.61 -4.40 -4.40 -4.38 -4.50 -4.58 -4.42 -4.47 -4.57 -4.45 -4.31

EQW må

Cet abundance dex

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

10.0

-4.21

95.5 14.0 18.8 39.4

-3.90 -4.25 -4.45 -4.14

46.2

-4.25

49.4

-4.13

7.66 111.40 5.15 0.93 6.18 3.42 8.43 4.06 7.38 59.51 10.48 4.62 8.15 1.51 3.01 8.93 7.07 0.89 2.91 4.61 3.09 2.83 5.16 2.82 8.10 60.31 2.54 9.32 3.08 3.58 15.68 6.87 2.55 21.59 5.56 8.68 3.15 35.87 3.86 2.01 1.94 28.06 2.77 0.76 5.00 18.05 62.50 5.83 30.70 18.18 4.70 4.47 8.27 1.14 3.10

91.7

-4.47

8.1

-4.63

13.0 8.4

-4.59 -4.58

9.4 2.6

-4.48 -4.65

3.4 2.3 11.2 15.3

-4.42 -4.61 -4.62 -4.85

39.3

-4.66

1.2 4.2 1.4 48.3 17.0 21.0 9.3

-4.60 -4.79 -4.33 -4.52 -4.55 -4.65 -4.62

28.3 96.6 50.7 27.3

-4.13 -3.96 -4.24 -4.18


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 23 Table 9. continued.

Element Wavelength å 5314.0342 5315.0859 5315.5630 5315.9961 5316.2251 5316.6152 5316.7842 5318.0571 5318.7500 5322.2339 5325.5532 5339.5850 5344.0952 5355.4189 5358.8760 5359.2461 5360.4858 5366.2070 5371.2749 5379.2441 5387.0630 5388.0210 5393.8472 5399.5640 5402.0591 5402.9072 5408.8110 5411.3750 5412.7461 5414.0732 5414.8501 5425.2568 5427.8262 5429.9878 5432.9668 5439.7070 5440.0679 5443.4492 5444.3872 5445.8071 5450.0991 5451.3179 5457.7300 5465.9312 5467.4551 5472.8550 5475.8291 5479.4009 5480.9551 5481.2759 5482.3081 5487.6190 5488.7822 5492.3989 5493.8330 5498.5762 5502.6709 5505.2559

excit eV 10.480 10.545 8.230 10.523 10.419 3.153 3.221 10.480 10.419 10.455 3.221 10.452 10.531 10.500 10.500 10.503 10.545 10.503 10.545 10.448 10.522 10.452 10.452 10.523 10.562 10.560 5.956 10.600 10.523 3.221 10.522 3.199 6.724 10.596 3.267 6.729 6.729 10.480 10.600 10.545 10.623 10.500 10.629 10.623 6.807 10.500 10.500 10.560 10.523 10.596 10.562 10.596 10.596 10.500 10.500 10.600 10.562 7.653

logg f -0.892 -0.418 -1.458 -1.129 0.340 -2.014 -2.783 -0.226 -0.544 -0.547 -3.324 0.517 -0.765 -0.500 -0.609 -0.675 -0.585 -0.196 -0.696 -0.988 0.499 -0.694 -0.246 -0.747 0.469 -0.591 -2.656 -0.047 -0.596 -3.645 -0.324 -3.390 -1.581 0.427 -3.527 -2.382 -2.739 -0.595 -0.170 -0.109 -0.093 -0.649 -0.138 0.348 -2.611 -0.715 -0.080 -0.353 -0.486 -0.278 0.413 0.288 -0.397 -0.097 0.259 -0.336 -0.192 -2.172

HD 145788 EQW abundance må dex

30.8

-4.06

9.4 23.7 11.9 22.1 14.2 21.5 9.1 40.9 23.2

-4.11 -4.20 -3.95 -4.20 -3.90 -4.36 -3.98 -4.06 -4.24

EQW må 1.78 3.36 5.08 1.47 15.50 68.48 42.06 6.78 3.06 3.53 22.41 17.96 1.65 3.53 3.42 1.78 2.14 5.37 1.97 1.55 16.84 2.03 4.94 2.19 16.50 2.00 6.25 4.82 1.53 11.71 20.33 13.88 13.45 13.86 3.68 1.83 3.04 4.80 5.47 2.31 2.38 4.90 13.72 2.32 2.90 5.53 3.14 1.98 2.48 14.16 2.90 6.84 11.67 3.10 5.62 2.32

21 Peg abundance dex -4.41 -4.56 -4.46 -4.24 -4.52 -4.45 -4.42 -4.43 -4.54 -4.44 -4.43 -4.57 -4.55 -4.46 -4.37 -4.61 -4.59 -4.56 -4.52 -4.38 -4.56 -4.55 -4.57 -4.42 -4.52 -4.60 -4.33 -4.71 -4.74 -4.49 -4.43 -4.56 -4.59 -4.49 -4.45 -4.41 -4.43 -4.58 -4.60 -5.00 -4.49 -4.58 -4.48 -4.39 -4.33 -4.64 -4.62 -4.72 -4.79 -4.55 -4.59 -4.50 -4.55 -4.61 -4.48 -4.39

EQW må

Cet abundance dex

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

7.5 3.4 13.9 20.5 3.5

-4.48 -4.57 -4.38 -4.65 -4.65

3.4 22.1 7.2

-4.88 -4.54 -4.48

4.5 5.3

-4.33 -4.77

11.2 12.4 15.6

-4.44 -4.54 -4.66

3.5 6.1 3.8 2.7 5.0 3.0 6.0 3.4 16.8 14.8 2.8 12.4 3.3

-4.47 -4.56 -4.87 -4.53 -4.66 -4.41 -4.72 -4.68 -4.63 -4.55 -4.71 -4.67 -4.64

15.2

-3.88

18.9

-4.14

21.3 11.1 18.1 10.6

-4.16 -4.22 -4.14 -4.00


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 24 Table 9. continued.

Element Wavelength å 5506.1948 5507.0718 5511.0830 5525.1250 5529.0532 5529.9321 5530.3408 5532.0879 5532.8638 5534.8472 5543.9551 5544.1958 5544.7632 5549.0010 5558.2871 5559.7949 5561.1499 5563.3989 5567.8418 5575.0210 5579.9248 5581.6279 5583.9438 5587.1138 5588.0298 5588.2202 5591.3682 5627.4971 5629.8711 5637.3428 5643.8799 5645.3921 5646.2261 5648.9038 5655.3570 5668.6362 5690.9941 5703.2539 5716.5898 5726.5570 5746.5752 5747.8838 5751.4878 5780.1279 5780.3359 5783.6299 5784.4482 5802.7681 5807.0181 5823.1548 5830.3408 5835.4922 5838.9888 5842.2900 5854.1919 5871.7988 5885.0151 5902.8252

excit eV 10.522 10.523 10.629 3.267 10.523 6.729 10.545 10.523 6.730 3.245 5.571 10.623 10.522 10.523 10.545 10.623 10.522 10.629 6.730 10.629 10.629 10.600 10.545 6.729 10.600 10.596 3.267 3.387 10.562 7.706 7.653 10.562 10.623 10.562 10.629 10.629 10.678 10.600 7.773 10.714 10.629 5.571 10.629 10.678 10.714 10.714 10.737 10.678 7.806 5.569 10.751 5.911 10.845 10.737 10.737 10.829 10.751 10.714

logg f 0.859 -0.056 -0.378 -4.102 -0.258 -1.813 -0.654 -0.099 -2.720 -2.865 -3.127 -0.232 0.139 -0.186 -0.781 -0.752 -0.581 -0.552 -1.866 -0.836 -0.503 -0.521 -0.826 -1.992 -0.740 0.163 -4.590 -4.188 -0.891 -2.253 -1.346 0.193 -0.584 -0.165 -0.550 -0.901 -0.175 -0.744 -2.075 -0.016 -0.402 -2.949 -0.611 0.421 -0.326 0.365 0.145 -0.452 -2.192 -2.987 -0.199 -2.702 -0.601 -0.328 -0.113 -0.277 0.298 0.416

HD 145788 EQW abundance må dex 35.9 -4.19 10.1 9.3 15.6 -4.17 -4.02 -4.24

64.3 15.9

-4.04 -4.10

EQW må 26.03 6.20 2.99 4.80 5.98 9.09 2.08 4.67 2.23 37.14 3.15 4.79 10.57 2.62 1.40 2.74 2.44 8.50 1.26 2.38 2.78 1.69 7.19 1.91 7.89 1.98 4.02 1.77 1.56 7.20 1.86 4.73 1.31 4.65 1.80 1.08 2.33 3.79 1.53 10.84 2.94 8.43 5.71 3.10 1.19 3.72 3.40 5.03 2.05 3.76 3.54 3.99 6.78 10.79

21 Peg abundance dex -4.54 -4.58 -4.57 -4.44 -4.40 -4.56 -4.51 -4.69 -4.33 -4.44 -4.37 -4.48 -4.48 -4.27 -4.55 -4.45 -4.49 -4.54 -4.51 -4.55 -4.47 -4.42 -4.49 -4.42 -4.62 -4.36 -4.37 -4.32 -4.45 -4.65 -4.57 -4.57 -4.40 -4.50 -4.42 -4.74 -4.62 -4.45 -4.61 -4.61 -4.54 -4.69 -4.67 -4.41 -4.56 -4.41 -4.57 -4.37 -4.36 -4.40 -4.64 -4.37 -4.71 -4.57

EQW må 29.7 5.4 6.4 5.1

Cet abundance dex -4.64 -4.78 -4.49 -4.76

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

13.5 6.7

-4.47 -4.54

16.2

-4.16

6.9

-4.53

13.0 12.3 12.2 16.1 17.2 11.5

-4.15 -4.23 -3.91 -4.15 -4.07 -3.94

6.2 7.2 2.6 5.3 6.5

-4.70 -4.49 -4.57 -4.58 -4.61

10.0

-4.44

10.1 5.3

-4.60 -4.86

11.0

-3.97

2.9 3.8 4.4 4.1 14.0

-4.60 -4.65 -4.68 -4.51 -4.62


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 25 Table 9. continued.

Element Wavelength å 5913.7710 5948.4268 5952.5098 5955.6982 5961.7051 5965.6221 5976.6821 5988.0112 5991.3760 6049.4448 6069.6748 6071.4258 6084.1108 6103.4961 6113.3218 6129.7031 6147.7412 6149.2578 6175.1460 6179.3838 6199.1812 6233.5342 6238.3921 6239.9531 6247.3501 6247.5571 6248.8979 6269.9668 6291.8301 6305.2959 6317.3940 6317.9829 6331.9541 6357.1621 6362.4741 6369.4619 6375.7920 6383.7222 6385.4512 6386.7129 6407.2510 6416.9189 6425.7202 6432.6802 6433.8140 6442.9551 6446.4102 6456.3828 6482.2041 6506.3330 6516.0801 6621.9790 6627.2612 6820.0390 6855.6460 6862.5259 6922.0298 6952.6719

excit eV 7.845 10.737 5.956 10.737 10.678 10.678 10.678 10.751 3.153 10.714 10.714 10.714 3.199 6.217 3.221 3.199 3.889 3.889 6.223 5.569 5.569 5.484 3.889 3.889 6.209 3.892 5.511 3.245 10.930 6.219 6.223 5.511 6.217 10.909 10.909 2.891 10.934 5.553 5.553 6.803 3.889 3.892 11.017 2.891 6.219 5.549 6.223 3.903 6.219 5.589 2.891 11.017 7.274 11.237 11.289 11.196 11.149 11.291

logg f -2.071 -0.201 -2.388 0.252 0.675 0.068 -0.335 -0.406 -3.647 -0.374 -0.502 -0.248 -3.881 -2.325 -4.230 -4.740 -2.827 -2.841 -2.086 -2.797 -2.941 -2.832 -2.754 -3.573 -2.172 -2.435 -2.784 -4.500 0.348 -2.094 -2.435 -2.155 -2.071 0.235 -0.492 -4.231 -0.011 -2.414 -2.715 -2.643 -3.854 -2.650 -0.007 -3.520 -2.738 -2.671 -2.082 -2.185 -1.853 -2.899 -3.432 0.055 -1.768 -0.280 -0.435 0.575 0.952 0.476

HD 145788 EQW abundance må dex

21.7

-4.27

19.6 12.0 16.1 45.9 44.4 19.0

-4.06 -4.10 -3.80 -4.06 -4.07 -4.07

EQW må 2.21 4.31 7.96 6.82 16.71 5.01 4.31 2.01 11.57 2.45 2.38 3.65 7.87 6.98 3.41 1.59 24.57 23.98 6.87 3.04 3.89 26.12 7.41 7.31 38.89 5.95 1.94 5.02 9.65 6.68 17.87 9.86 5.04 1.83 5.22 3.05 10.37 6.48 1.75 3.38 21.71 2.90 14.28 2.61 6.60 9.22 47.20 13.72 2.70 25.21 2.26 5.26 5.01 11.80 4.35

21 Peg abundance dex -4.36 -4.44 -4.42 -4.65 -4.56 -4.66 -4.33 -4.58 -4.50 -4.53 -4.41 -4.46 -4.44 -4.40 -4.47 -4.32 -4.43 -4.43 -4.28 -4.52 -4.56 -4.45 -4.38 -4.51 -4.40 -4.39 -4.44 -4.75 -4.44 -4.28 -4.40 -4.45 -4.63 -4.39 -4.44 -4.62 -4.43 -4.38 -4.39 -4.46 -4.67 -4.59 -4.63 -4.42 -4.41 -4.46 -4.38 -4.46 -4.58 -4.34 -4.74 -4.46 -4.73 -4.62 -4.63

EQW må 6.6 21.1 9.5

Cet abundance dex -4.39 -4.62 -4.49

Ref logg f RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU RU

3.8 4.2 3.4 5.7 15.1 14.4 5.8 17.3

-4.35 -4.56 -4.50 -4.39 -4.48 -4.49 -4.21 -4.46

46.7 72.7 15.6 15.8 22.9 22.9 20.6

-4.11 -3.82 -3.90 -4.16 -3.96 -3.78 -4.07

3.7

-4.45

9.4

-4.38

2.7

-4.87

26.2 16.2 16.5 83.3 23.4

-4.24 -3.96 -4.14 -3.80 -4.15

12.8 3.8 7.3 3.8 7.0

-4.50 -4.67 -4.46 -4.53 -4.50

6.6 2.9 1.0 7.0

-4.35 -4.38 -4.69 -4.82


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 26 Table 9. continued.

Element Wavelength å 6954.8599 7135.0239 7224.4870 7320.6540 7449.3350 7462.4070 7513.1620 7533.3670 7981.8980 Fe III 4371.3370 4419.5962 4431.0190 5156.1108 Co II 4160.6730 4413.9070 4569.2500 4660.6560 Ni I 3858.2920 4401.5381 4648.6460 4714.4082 4756.5098 4786.5308 4904.4072 4980.1660 4984.1118 5035.3569 5080.5278 5081.1069 5476.8999 Ni II 3849.5540 4015.4741 4067.0310 4187.8490 4192.0649 4244.7788 4279.2158 4509.2720 4665.5479 4679.1592 4791.2290 4898.9448 4932.6580 4974.1300 4981.0000 5003.4141 5014.6201 5021.4712 5040.3452 5042.4302 5052.9971 5058.3760 5059.2002 5066.3281

excit eV 11.255 6.209 3.889 3.892 3.889 3.892 9.654 3.903 9.654 8.241 8.241 8.248 8.641 3.408 5.046 3.406 3.361 0.423 3.193 3.420 3.380 3.480 3.420 3.542 3.606 3.796 3.635 3.655 3.847 1.826 4.032 4.032 4.029 4.029 4.032 4.032 6.616 6.764 6.856 6.989 4.029 12.208 12.220 13.976 12.253 12.540 12.554 12.437 12.271 12.568 12.294 12.567 12.291 6.329

logg f -0.171 -2.602 -3.358 -3.336 -3.448 -2.871 0.294 -3.600 -0.560 -3.031 -2.210 -2.575 -1.995 -1.828 -1.853 -2.399 -2.345 -0.94(02) 0.04(03) -0.10(02) 0.26(02) -0.27(02) -0.16(03) -0.17(18) 0.07(02) 0.226 0.29(02) 0.33(02) 0.30(18) -0.89(07) -1.878 -2.424 -1.835 -2.676 -3.055 -3.109 -1.997 -2.185 -1.821 -1.748 -3.841 0.435 0.185 0.631 0.285 0.702 0.281 0.920 0.574 0.574 0.366 0.846 0.539 -1.799

HD 145788 EQW abundance må dex

EQW må 1.96 2.32

21 Peg abundance dex -4.40 -4.55

EQW må 10.5 6.8 6.5 4.1 12.0 10.5 3.3 2.8 1.6 8.1 6.0 6.7 2.5

Cet abundance dex -4.98 -4.31 -4.35 -4.45 -4.49 -4.98 -4.38 -4.74 -4.61 -4.57 -4.38 -4.53 -6.93

Ref logg f RU RU RU RU RU RU RU RU RU Kur06 Kur06 Kur06 Kur06 RPU RPU RPU RPU BBPL WLa WLa WLa WLa WLa FMW WLa K88 WLa WLa FMW FMW

2.79 1.09 1.54 1.55 1.55 1.70 S 19.5 -5.32 -5.25 2.60 1.35 2.70 0.93 1.20 0.88 1.94 2.98 2.40 2.03

-4.56 -4.67 -4.57 -6.54 -6.83 -6.86 -5.69 -5.72 -5.79 -5.68 -5.71 -5.77 -5.67 -5.63 -5.76 -5.70

8.0 8.4 18.2 10.6 10.4 17.0 S

-5.40 -5.42 -5.18 -5.49 -5.36 -5.10 -5.22

18.96 38.53 16.1 15.7 -5.16 -5.12 7.13 5.90 4.53 2.51 4.29 4.07 1.71 1.93

-5.68 -5.67 -5.58 -5.62 -5.57 -5.56 -5.61 -5.64 -5.42 -5.67

10.1 24.9 8.0 4.0 3.1 3.4 3.8 1.3 2.4 2.5 1.5 2.9

-5.88 -5.94 -5.74 -5.69 -5.72 -5.72 -5.68 -6.09 -5.55 -5.25 -5.85 -5.84

10.4

-5.07

2.33 0.85 3.01 3.08 2.82 1.98 3.76 2.69 6.00

-5.67 -5.71 -5.81 -5.52 -5.43 -5.51 -5.55 -5.54 -5.67

2.8 2.2 2.4 3.1 5.2

-5.71 -5.73 -6.05 -5.74 -5.75

K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88 K88


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 27 Table 9. continued.

Element Wavelength å 5105.1382 5139.9219 5424.5718 6124.8892 Zn I 4810.5278 Sr II 4077.7090 4215.5190 Y II 4177.5290 4374.9350 4900.1200 5205.7240 5662.9250 Zr II 4029.6841 4149.2168 4208.9771 4443.0078 Ba II 4554.0288 4934.0762 6141.7129 6496.8970 Nd III 4927.4877 5203.9236 5294.1133

excit eV 12.317 12.329 6.473 6.392 4.078 0.000 0.000 0.409 0.409 1.033 1.033 1.944 0.713 0.802 0.713 1.486 0.000 0.000 0.704 0.604 0.461 0.141 0.000

logg f 0.392 0.532 -1.988 -2.048 -0.237 0.167 -0.145 -0.16(02) 0.16(01) -0.09(03) -0.34(03) 0.16(02) -0.78(02) -0.04(02) -0.51(02) -0.42(03) 0.170(028) -0.150(030) -0.076(044) -0.377(040) -0.800 -0.660 -0.690

HD 145788 EQW abundance må dex

EQW må 1.33 2.09 2.72 3.05 2.56

21 Peg abundance dex -5.70 -5.61 -5.74 -5.58 -6.86 -9.09 -9.10 -9.71 -9.96 -9.61 -9.76 -9.21 -9.59 -9.82 -9.28 -9.24 -9.19 -9.12 -10.16 -10.07 -10.03

EQW må

Cet abundance dex

Ref logg f K88 K88 K88 K88 Wb

S 78.9

-8.45 -8.45

S 27.59 3.13 3.74 2.22

S 3.0

-9.14 -9.16

Wa Wa HL HL HL HL HL LNAJ LNAJ LNAJ LNAJ MW MW MW MW RRKB RRKB RRKB

13.5 18.4 28.2 10.5 48.1 34.7 28.9

-9.06 0.87 -8.38 -8.81 -8.96 -9.00 -9.05 -8.84 1.43 2.90 0.68 1.09 15.25 5.14 3.55 0.60 1.50 1.82

Wavelengths and excitation potentials are taken from the VALD database. The adopted logg f values are taken from different sources which are listed in the last column. Errors in logg f values if available are given in paranthesis. "S" means that the line abundance was determined by fitting an observed line profile and not with equivalent width. AJPP81 - Artru et al. (1981); BBPL - Blackwell et al. (1989); ´ BGF - Biemont et al. (1989); BHN - Bizzari et al. (1993); BK - Bard & Kock (1994); BKK - Bard et al. (1991); BL - O'Brian & Lawler (1991); BM - Bengston & Miller (1970); BWL - O'Brian et al. (1991); FMW - Fuhr et al. (1988); H - Hibbert (1988); HL - Hannaford et al. (1982); K88 - Kurucz (1988), and KurXX - Kurucz calculations of the XX year; KG - Kling & Griesmann (2000); KP - Kurucz & Peytremann (1975); KP08 - Kelleher & Podobedova (2008); KSG - Kling et al. (2001); L - Lilly (1976); LD - Lawler & Dakin (1989); LNAJ - Ljung et al. (2006); MFW - Martin et al. (1988); MRB - Miller et al. (1971); Mult - estimated from multiplet table intensity;


Fossati et al.: The abundance analysis for normal A- and B-type stars., Online Material p 28

MW - Miles & Wiese (1969); MWRB - Miller et al. (1974); NIST08 - Ralchenko et al. (2008); PTP - Pickering et al. (2001); RHL - Ryabchikova et al. (1994); RPU - Raassen et al. (1998); RRKB - Ryabchikova et al. (2006); RU - Raassen & Uylings database (ftp://ftp.wins.uva.nl/pub/orth); S - Smith (1988); SG - Smith & Gallagher (1966); SN - Smith & O'Neil (1975); SR - Smith & Raggett (1981); SW - Smith & Wiese (1971); T - Theodosiou (1989); TB - TopBase, Seaton et al. (1994); Wa - Warner (1968a); Wb - Warner (1968b); WF - Wiese & Fuhr (2007); WLa - Wickliffe & Lawler (1997); WSG - Wiese et al. (1966); WSM - Wiese et al. (1969).