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Astronomy & Astrophysics manuscript no. teff (DOI: will be inserted by hand later) February 2, 2008 High precision effective temperatures for 181 F–K dwarfs from line-depth ratios ⋆ arXiv:astro-ph/0308429v1 25 Aug 2003 V.V. Kovtyukh1 , C. Soubiran2 , S.I. Belik1 , and N.I. Gorlova3 1 2 3 Astronomical Observatory of Odessa National University and Isaac Newton Institute of Chile, Shevchenko Park, 65014, Odessa, Ukraine email: [email protected] Observatoire de Bordeaux, CNRS UMR 5804, BP 89, 33270 Floirac, France Steward Observatory, The University of Arizona, Tucson, AZ, USA 85721 Received 6 May 2003; accepted Abstract. Line depth ratios measured on high resolution (R=42 000), high S/N echelle spectra are used for the determination of precise effective temperatures of 181 F,G,K Main Sequence stars with about solar metallicity (–0.5 < [Fe/H] < +0.5). A set of 105 relations is obtained which rely Teff on ratios of the strengths of lines with high and low excitation potentials, calibrated against previously published precise (one per cent) temperature estimates. The application range of the calibrations is 4000–6150 K (F8V–K7V). The internal error of a single calibration is less than 100 K, while the combination of all calibrations for a spectrum of S/N =100 reduces uncertainty to only 5–10 K, and for S/N =200 or higher – to better than 5 K. The zero point of the temperature scale is directly defined from reflection spectra of the Sun with an uncertainty about 1 K. The application of this method to investigation of the planet host stars properties is discussed. Key words. Stars: fundamental parameters – stars:temperatures – stars:dwarfs – planetary systems 1. Introduction The determination of accurate effective temperatures is a necessary prerequisite for detailed abundance analysis. In this paper we focus on dwarfs with solar metallicity (–0.5 < [Fe/H] < +0.5) to contribute to the very active research field concerning the fundamental parameters of stars with planets. High precision temperatures of such stars might help to resolve two outstanding questions in the extra-solar planetary search. Namely, to get a definite confirmation of the metal richness of the stars that harbor planets, and secondly, perhaps to rule out some lowmass planetary candidates by detecting subtle variations in the host’s temperature due to star-spots. The numerous studies of the large fraction of the known extra-solar planet hosts (∼80 out of ∼100 known systems) have revealed their larger than average metal richness (Gonzalez 1997; Fuhrmann, Pfeiffer & Bernkopf 1998; Gonzalez et al. 2001 and references therein; Takeda et al 2001; Santos et al. 2003 and references therein). The reliability of this Send offprint requests to: V.V. Kovtyukh, e-mail: [email protected] ⋆ Based on spectra collected with the ELODIE spectrograph at the 1.93-m telescope of the Observatoire de Haute Provence (France). result depends mainly on the accuracy of the model atmosphere parameters, with effective temperature (Teff ) being the most important one. The direct method to determine the effective temperature of a star relies on the measurement of its angular diameter and bolometric flux. In practice certain limitations restrict the use of this fundamental method to very few dwarfs. Other methods of temperature determination have errors of the order 50–150 K, which translates into the [Fe/H] error of ∼0.1 dex or larger. The only technique capable so far of increasing this precision by one order of magnitude, is the one employing ratios of lines with different excitation potentials χ. As is well known, the lines of low and high χ respond differently to the change in Teff . Therefore, the ratio of their depths r = Rλ1 /Rλ2 (or equivalent widths, EW) should be very sensitive temperature indicator. The big advantage of using line-depth ratios is the independence on the interstellar reddening, spectral resolution, rotational and microturbulence broadening. The reader is referred to Gray (1989, 1994) and Gray & Johanson (1991) to learn more about the history and justification of the line ratio method. Applying this method to the Main-Sequence (MS) stars, they achieved precision as high as 10 K. The most recent works on the subject are by 2 Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs Caccin, Penza & Gomez (2002) who discuss the possible weak points of this technique for the case of dwarfs (see below), and the fundamental contribution by Strassmeier & Schordan (2000) who report 12 temperature calibrations for giants with an error of only 33 K. So far however the line-ratio method has not been fully utilized for the purposes other than just temperature estimation by itself. One of few applications is the chemical abundance analysis of supergiants, where it has proved the anticipated high efficiency and accuracy. Thus, Kovtyukh & Gorlova (2000, hereafter Paper I) using high-dispersion spectra, established 37 calibrations for the temperature determination in supergiants (a further study increased this number to 55 calibrations). Based on this technique, in the series of 3 papers Andrievsky et al. (2002 and references therein) derived temperatures for 116 Cepheids (from about 260 spectra) at a wide range of galactocentric distances (Rg =5–15 kpc) with a typical error 5–20 K. The high precision of this new method of temperature determination allowed them to uncover the fine structure in the Galactic abundance gradients for many elements. Even for the most distant and faint objects (V ≃ 13–14 mag) the mean error in Teff was no larger than 50–100 K, with maximum of 200 K for spectra with lowest S/N (=40–50). Another example concerns T Tau stars. For young stars, uncertainties in reddening due to variable circumstellar extinction invalidate the photometric color method of effective temperature determination. Using 5 ratios of FeI and VI lines calibrated against 13 spectral standards, Padgett (1996) determined the effective temperature of 30 T Tau stars with a 1 σ uncertainty lower than 200 K. The intent of this paper is to improve this technique, based on our experience of applying it to supergiants (Paper I and following publications), and expand it to the MS stars. The wide spectral range of ELODIE echelle spectra allowed to select many unblended lines of low and high excitation potentials thus improving the internal consistency of the method, whereas the large intersection between the ELODIE database and published catalogues of effective temperatures allowed to take care of systematic effects. We obtained a median precision of 6 K on Teff derived for an individual star. The zero-point of the scale was directly adjusted to the Sun, based on 11 solar reflection spectra taken with ELODIE, leading to the uncertainty in the zero-point of about 1 K. Temperature determined by the line ratio method may now be considered as one of the few fundamental stellar parameters that have been measured with internal precision of better than 0.2%. 2. Observations and temperature calibrations The investigated spectra are part of the library collected with the ELODIE spectrometer on the 1.93-m telescope at the Haute-Provence Observatory (Soubiran et al. 1998, Prugniel & Soubiran 2001). The spectral range is 4400– 6800 ÅÅ and the resolution is R=42000. The initial data reduction is described in Katz et al. (1998). All the spec- tra are parametrized in terms of Teff , logg, [Fe/H], either collected from the literature or estimated with the automated procedure TGMET (Katz et al. 1998). This allowed us to select a sample of spectra of FGK dwarfs in the metallicity range –0.5< [Fe/H] < +0.5. Accurate Hipparcos parallaxes are available for all of the stars of interest enabling to determine their absolute magnitudes MV that range between 2.945 (HD81809, G2V) and 8.228 (HD201092, K7V). All the selected spectra have a signal to noise ratio greater than 100. Further processing of spectra (continuum placement, measuring equivalent widths, etc.) was carried out by us using the DECH20 software (Galazutdinov 1992). Equivalent widths EWs and depths Rλ of lines were measured manually by means of a Gaussian fitting. The Gaussian height was then a measure of the line depth. This method produces line depths values that agree nicely with the parabola technique adopted in Gray (1994). We refer the reader to Gray (1994, and references therein), Strassmeier & Schordan (2000) for a detailed analysis of error statistics. Following Caccin’s, Penza & Gomez (2002) results, where a careful analysis of the anticipated problems for the Solar-type stars has been carried out, we did not use ion lines and high-ionization elements (like C, N, O) due to their strong sensitivity to gravity. Gray (1994) showed that the ratio of lines VI 6251.82 and FeI 6252.55 depends strongly on metallicity. The reason is that the strong lines like FeI 6252.55 (Rλ =0.52 for the Sun) are already in the dumping regime, where the linearity of EW on abundance breaks down. In addition, as was shown in the careful numerical simulations by Stift & Strassmeier (1995), this ratio (of 6251.82 and 6252.55 lines) is also sensitive to rotational broadening. Significant effects were found for vsini as small as 0–6 km s−1 (for solar-like stars). We therefore avoided to use strong lines in our calibrations. Indeed, Gray (1994) concluded that, as expected, the weak-line ratios are free from the effects of metallicity. As to the effect of rotation, we should note that all objects in our sample are old Main Sequence stars with slow to negligible rotation (vsini<15 km s−1 ), which is comparable to the instrumental broadening. Thus, we initially selected about 600 pairs of 256 unblended SiI, TiI, VI, CrI, FeI, NiI lines with high and low excitation potentials within the wavelength interval 5300– 6800 ÅÅ. These lines have been selected according to the following criteria: (1) the excitation potentials of the lines in pair must differ as much as possible; (2) the lines must be close in wavelength; it turned out though that calibrations based on widely spaced lines (including from different orders) show same small dispersion as the closely spaced lines. Therefore, we retained all pairs with difference in wavelength up to 70 Å(λ2 − λ1 <70 Å); (3) the lines must be weak enough to eliminate a possible dependence on microturbulence, rotation and metallicity; Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs 3 1.0 0.6 6080 6081 6082 6083 6084 6085 6086 6087 o 6088 6089 6090 6091 6091.92 SiI 5.87 6091.18 TiI 2.27 6090.21 VI 1.08 6086.29 NiI 4.27 6085.27 FeI 2.76 0.7 6089.57 FeI 5.02 0.8 6081.44 VI 1.05 Relative Flux 0.9 6092 6093 Lambda, A Fig. 1. Comparison spectra for two stars: solid line – a planet-host star HD 217014 (51 Peg), and dotted line – a non-planet star HD 5294. Within the limits of the errors, both stars have identical temperatures (5778 and 5779 K, respectively), but different metallicities. Spectral lines used in temperature calibrations are identified at the bottom with their wavelength, element, and lower excitation potentials χ in eV. We used ratios 6081.44/6089.57, 6085.27/6086.29, 6085.27/6155.14, 6089.57/6126.22, 6090.21/6091.92, 6090.21/6102.18, 6091.92/6111.65 and others. (4) the lines must be situated in the spectral regions free from telluric absorption. The next step was to choose the initial temperatures for interpolation. This is a very important procedure since it affects the accuracy of the final temperature scale, namely, the run of the systematic error with Teff (Fig. 3). There is an extended literature on MS stars temperatures. For 45 stars from our sample (see Table 1) we based the initial temperature estimates on the following 3 papers: Alonso, Arribas & Martinez-Roger (1996, hereafter AAMR96), Blackwell & Lynas–Gray (1998, hereafter BLG98) and DiBenedetto (1998, hereafter DB98). In these works the temperatures have been determined for a large fraction of stars from our sample with a precision of about 1%. AAMR96 used the Infrared Flux Method (IRFM) to determine Teff for 475 dwarfs and subdwarfs with a mean accuracy of about 1.5% (i.e., 75–90 K). BLG98 also have determined temperatures for 420 stars with spectral types between A0 and K3 by using IRFM and achieved an accuracy of 0.9%. DB98 derived Teff for 537 dwarfs and giants by the empirical method of surface brightness and Johnson broadband (V − K) color, the accuracy claimed is ±1%. Whenever 2 or 3 estimates were available for a given star, we averaged them with equal weights. These temperatures served as the initial approximations for our calibrations. First, for the above mentioned 45 stars with previously accurately determined Teff we plotted each line ratio against Teff , and retained only those pairs of lines that showed unambiguous and tight correlation. We experimented with a total of nearly 600 line ratios but adopted only the 105 best - the ones showing the least scatter. These 105 calibrations consist of 92 lines, 45 with low ( χ <2.77 eV ) and 47 with high ( χ >4.08 eV ) excita- tion potentials. Judging by the small scatter in our final calibrations (Fig.2) and Teff , the selected combinations are only weakly sensitive to effects like rotation, metallicity and microturbulence. This confidence is reinforced by the fact that the employed lines belong to the wide range of chemical elements, intensity and atomic parameters, therefore one can expect the mutual cancellation of the opposite effects. Each relationship was then fitted with a simple analytical function. Often calibrations show breaks which can not be adequately described even by a 5th-order polynomial function (see Fig. 2). Therefore, we employed other functions as well, like Hoerl function (Teff =abr ∗ rc , where r = Rλ1 /Rλ2 , a, b, c – constants), modified Hoerl (Teff =ab1/r rc ), power low (Teff =arb ), exponential (Teff =abr ) and logarithmic (Teff =a+b ln(r)) functions. For each calibration we selected function that produced the least square deviation. As a result, we managed to accurately approximate the observed relationships with a small set of analytic expressions. This first step allowed to select 105 combinations, with an rms of the fit lower than 130 K, the median rms being 93 K. Using these initial rough calibrations, for each of the 181 target stars we derived a set of temperatures (70–100 values, depending on the number of line ratios used), averaged them with equal weights, and plotted these mean Teff (with errors of only 10–20 K) versus line ratios again, thus determining the preliminary calibrations (for which the zero-point had yet to be adjusted). We would like to point out that the precision of our calibrations varies with temperature. In particular, at high Teff the lines with low χ become very weak causing line depth measurement to be highly uncertain. Therefore, for 4 Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs each calibration we determined the optimum temperature range where the maximum accuracy is attained (no worse than 100 K), so that for a given star only a subset of calibrations can be applied. 6500 6000 Teff, K Teff, K (other study) 5500 Teff, K 6000 6000 5500 5500 SiI 5690.43 / VI 5703.59 5000 5000 4500 4000 3500 3500 SiI 5701.11 / VI 5727.05 5000 4000 4500 5000 5500 6000 6500 Teff, K (present paper) 4500 4500 0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 6000 6000 TiI 6126.22 / SiI 6155.14 VI 5727.05 / SiI 5772.15 Fig. 3. Comparison between the temperatures derived in the present work and those derived by AAMR96 – squares, BLG98 – circles, and DB98 – triangles. The dashed line represents the linear fit to the data, and the solid line represents the one-to-one correlation 5500 35 5000 5000 4500 0.0 1.0 2.0 3.0 4.0 5.0 6000 4000 0.0 1.0 2.0 3.0 4.0 5.0 Teff standart error, K 30 25 20 15 10 5 6000 0 4000 5500 5000 5000 NiI 6223.99 / VI 6251.82 4500 0.0 4500 5000 5500 6000 Teff, K 5500 Fig. 4. Standard error of the mean versus effective temperature averaged over all available line ratios. SiI 6243.81 / TiI 6261.10 4500 0.5 1.0 1.5 2.0 2.5 3.0 R1/R2 3.5 0.0 1.0 2.0 3.0 4.0 5.0 R1/R2 Fig. 2. Our final calibrations of temperature versus line depth ratios r=R1/R2. The temperatures are shown as the average value derived from all calibrations available for a given star. The errors in temperature are less than the symbol size. The typical error in line ratio is 0.02–0.05. Position of the Sun is marked by the standard symbol. What are the main sources of random errors in the line ratio method? The measurement errors in line depths are mainly caused by the continuum placement uncertainty and by the Gaussian approximation of the line profile. In addition, the individual properties of the stars, such as metallicity, spots, rotation, convection, non-LTE effects, and binarity may also be responsible for the scatter observable in Fig. 2. The detailed analysis of these and other effects can be found in Paper I, Strassmeier & Schordan 2000 and in works by D.F.Gray. We estimate that the typical error in the line depth measurement r = Rλ1 /Rλ2 is 0.02-0.05, implying an error in temperature of about 20–50 K. The mean random error of a single calibration is 60–70 K (40–45 K in the most and 90–95 K in the least accurate cases). The use of ∼70–100 calibrations reduces the uncertainty to 5–7 K (for spectra with S/N =100–150). Better quality spectra (R >100,000, S/N > 400) should in principle allow the uncertainty of just 1–2 K. Clearly, time variation of the temperature for a given star should be readily detected by this method, since the main parameters that cause scatter due to star-to-star dissimilarities ( gravity, rotation, [Fe/H], convection, non-LTE effects etc.) are fixed for a given star. The temperature variation of several degrees in mildly active stars may be produced by the surface features and rotational modulation, as for example has been documented for the G8 dwarf ξ Bootis A (Toner & Gray 1988) and σ Dra (K0V, Gray et al. 1992). The next stage is to define the zero point of our temperature scale. Fortunately, for dwarfs (unlike for supergiants) a well-calibrated standard exists, the Sun. Using our preliminary calibrations and 11 independent solar Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs Table 1. RMS of a linear regression between Teff and Strömgren b − y using effective temperatures obtained by other authors and in this study with N common stars. author AAMR96 BLG98 DB98 EDV93 N 30 25 29 30 σothers K 102 71 113 63 σour K 78 65 87 29 spectra from the ELODIE library (reflection spectra of the Moon and asteroids), we obtained a mean value of 5733 ±0.9K for the Sun’s temperature. Considering Sun a normal star, we adjusted our calibrations by adding 44 K to account for the offset between the canonical Solar temperature of 5777 K and our estimate. The possible reasons for this small discrepancy are discussed below. Another point concerns the difference between the zero-point of our temperature scale and that of other authors. Comparing 30 common objects, we find that AAMR96 scale underestimates temperatures by 45 K near the solar value compared to ours, but apart from that, the deviations are random, no trend with Teff is present. The 45 K offset may arise from the various complications associated with observing the Sun as a star, and/or problems in the models used by AAMR96, like underestimation of convection in the grid of the model atmosphere flux developed by Kurucz. After correcting AAMR96 zero-point for the 45 K offset, the mean random error of their scale becomes 65 K (where we neglect the error of our own scale which is an order less). The temperatures of BLG98 are also in a good agreement with our estimates – except for a 48 K offset, no correlation of the difference with temperature is observed. The mean dispersion is 63 K (for 26 common stars), which is within the errors of BLG98 scale. Comparing with DB98: for the 29 star in common, their temperatures are on average 41 K below ours, and the mean error is ±53K. Thus, the temperatures derived in AAMR96, BLG98 and DB98 have good precision, though the absolute values are somewhat low relative to the Sun. The reason may be due to the difficulty of the photometric measurements of the Sun, as well as indicate some problems in the model atmosphere calculations employed. For example, the Sun’s temperatures derived in AAMR96 and DB98 are identical – 5763 K, which is below the nominal value of 5777 K. Besides, the mean temperatures of solar analogue stars (spectral types G2-G3, [Fe/H]≈0.0, and Sun being of G2.5 type) derived in these papers, are significantly below the solar value: 5720±54K (AAMR96, 3 stars), 5692±31K (BLG98, 11 stars) and 5702±46K (DB98, 7 stars). Our determination for the G2-G3 spectral types is 5787±14K, based on 12 stars. This demonstrates that a small error (0.8%) affects the zero point of the IRFM method, because 5 when applied to the Sun and the solar type stars, it returns inconsistent results. We also compared our estimates of Teff with photometrical temperatures. EDV93 derived temperatures of 189 nearby field F,G disk dwarfs using theoretical calibration of temperature versus Strömgren (b–y) photometry (see Table 1). The mean difference between Teff of Edvardsson et al.(1993) and ours is only −14K (σ=±67 K, based on 30 common stars). To compare our temperatures to Gray (1994), we used his calibration of (B-V)corr corrected for metallicity. Our scale is +11 K lower (σ=±61 K, 24 stars). Summarizing, we demonstrated that our temperature scale is in excellent agreement with the widely used photometric scales, while both the IRFM method and the method of surface brightness predict too low values for the temperature of the Sun and the solar type stars. Fig. 4 shows the sensitivity of our technique to temperature. Two outliers with errors greater than 20 K are the cold dwarfs HD28343 and HD201092, known as flaring stars. For other stars the internal errors range between 3 and 13 K, with a median of 6 K. 3. Results and Discussion Table 1 contains our final Teff determinations for 181 MS stars. Note that we added the 44 K correction to the initial calibrations in order to reproduce the standard 5777 K temperature of the Sun. For each star we report the mean Teff , number of the calibrations used (N ), and the standard error of the mean (σ). For comparison, we also provide Teff as determined in Edvardsson et al. (1993, hereafter EDV93), AAMR96, BLG98 and DB98. Absolute magnitudes MV have been computed from Hipparcos parallaxes and V magnitudes from the Tycho2 catalogue (Høg et al. 2000) transformed into Johnson system. (B − V ) are also from Tycho2. Planet harboring stars are marked with an asterisk. As one can see from Table 1, for the majority of stars we get an error which is smaller than 10 K. The consistency of the results derived from the ratios of lines representing different elements is very reassuring. It tells that our 105 calibrations are essentially independent from micro-turbulence, LTE departures, abundances, rotation and other individual properties of stars. We admit though that a small systematic error may exist for Teff below 5000 K where we had only few standard stars. As was already mentioned, for the first approximation we took accurate temperatures from AAMR96, BLG98 and DB98. The comparison of our final Teff with those derived by AAMR96, BLG98 and DB98 is shown in Fig. 3. As a test of the internal precision of our Teff we investigate the Teff – color relation with the Strömgren index b–y, using our determinations of Teff , and those obtained by other authors. The results are shown in Table 2 where the rms of the linear fit is given for each author’s determination, along with our estimate of Teff and using common stars. In each case the scatter of the color relation is significantly 6 Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs Table 2. Program stars. Asterisks indicate stars with planets. HD/BD HR 1562 1835 3765 4307 4614 5294 6715 8574 8648 9826 10145 10307 10476 10780 11007 13403 13507 13825 14374 15335 17674 17925 18803 19019 19308 19373 19994 22049 22484 23050 24053 24206 26923 28005 28099 28343 28447 29150 29310 29645 29697 30495 – 88 – 203 219 – – – – 458 – 483 493 511 523 – – – – 720 – 857 – – – 937 962 1084 1101 – – – 1322 – – – – – – 1489 – 1532 Name 9 Cet 18 Cet 24 Eta Cas 50 Ups And 107 Psc 13 Tri Iot Per 94 Cet 18 Eps Eri 10 Tau V774 Tau 58 Eri Teff this paper 5828 5790 5079 5889 5965 5779 5652 6028 5790 6074 5673 5881 5242 5407 5980 5724 5714 5705 5449 5937 5909 5225 5659 6063 5844 5963 6055 5084 6037 5929 5723 5633 5933 5980 5778 4284 5639 5733 5852 6009 4454 5820 N σ, K 97 68 87 91 69 86 97 61 59 44 96 94 69 95 84 91 91 96 77 84 58 87 95 56 95 75 56 84 60 80 93 94 77 87 85 20 93 89 89 57 40 91 5.8 5.5 4.7 5.0 6.4 6.6 6.7 6.7 7.2 13.1 4.2 4.0 3.2 4.0 7.4 7.0 5.4 5.5 4.6 6.6 8.7 5.0 3.5 7.2 5.4 5.1 10.0 5.9 3.6 9.0 3.7 4.8 5.9 6.1 5.2 20.3 6.3 5.4 7.7 5.8 11.4 5.7 Teff EDV93 improved when adopting our temperatures, though some residual dispersion is still present that can be attributed to the photometric errors, reddening and the intrinsic properties of stars (metallicity, gravity, binarity...) to which the color indices are known to be sensitive sensitive. The improvement is particulary spectacular in comparison with EDV93. This proves the high quality of our temperatures and the mediocrity of b–y as a temperature indicator. 5809 5946 Teff AAMR96 Teff BLG98 Teff DB98 5713 5774 5753 5817 5771 6212 6155 6136 5898 5874 5172 5223 5157 5585 5577 5588 5857 5996 5875 5921 5880 5981 5951 6104 5981 6028 5076 5998 5944 5940 5781 5775 Mv B–V 5.006 4.842 6.161 3.637 4.588 5.065 5.079 3.981 4.421 3.452 4.871 4.457 5.884 5.634 3.612 3.949 5.123 4.700 5.492 3.468 4.194 5.972 4.998 4.445 4.220 3.935 3.313 6.183 3.610 4.330 5.183 5.418 4.685 4.359 4.747 8.055 3.529 4.934 4.407 3.504 7.483 4.874 0.585 0.621 0.954 0.568 0.530 0.610 0.658 0.535 0.643 0.496 0.667 0.575 0.819 0.767 0.524 0.616 0.637 0.674 0.757 0.539 0.563 0.864 0.669 0.508 0.626 0.554 0.523 0.877 0.527 0.544 0.674 0.681 0.537 0.652 0.660 1.363 0.678 0.668 0.564 0.548 1.108 0.588 rem * * * * * within 3–13 K (median 6 K) for the major fraction of the sample, except for the two outliers. We demonstrated that the line ratio technique is capable of detecting variations in Teff of a given star as small as 1–5 K. This precision may be enough to detect star spots and Solar-type activity cycles. Of particular interest is the application of this method to testing ambiguous cases of low-mass planet detection, since planets do not cause temperature variations, unlike spots. 4. Conclusion The high-precision temperatures were derived for a set of 181 dwarfs, which may serve as temperature standards in the 4000–6150 K range. These temperatures are precise to The next step will be the adaptation of this method to a wider range of spectral types and for an automatic pipeline analysis of large spectral databases. Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs Table 1 (Continued) HD HR 30562 32147 34411 38858 39587 40616 41330 41593 42618 42807 43587 43947 45067 47309 50281 50554 51419 55575 58595 60408 61606 62613 64815 65874 68017 68638 70923 71148 72760 72905 73344 75318 1536 1614 1729 2007 2047 – 2141 – – 2208 2251 – 2313 – – – – 2721 – – – 2997 – – – – – 3309 – 3391 – – Name 15 Lam Aur 54 Chi1 Ori 3 Pi1 UMa Teff this paper 5859 4945 5890 5776 5955 5881 5904 5312 5775 5737 5927 6001 6058 5791 4712 5977 5746 5949 5707 5463 4956 5541 5864 5936 5651 5430 5986 5850 5349 5884 6060 5450 N σ, K 87 65 88 81 71 89 77 92 96 81 81 82 61 95 56 77 94 65 87 97 83 90 88 85 100 90 82 88 91 79 37 78 6.8 8.7 4.3 6.7 6.1 10.0 5.5 3.3 6.6 5.2 4.4 7.1 4.6 3.9 8.5 5.8 8.3 6.6 8.3 4.7 4.6 6.4 8.3 4.7 9.0 6.3 4.5 5.1 3.8 6.8 6.8 5.8 7 Teff EDV93 5886 Teff AAMR96 5822 Teff BLG98 5843 Teff DB98 5871 5889 5847 5848 5669 5859 5697 5953 5917 5945 5963 Acknowledgements. V.K. wants to thank the staff of Observatoire de Bordeaux for the kind hospitality during his stay there. The authors are also grateful to the anonymous referee for the careful reading of the manuscript and the numerous important remarks that helped to improve the paper. 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Vienna, p. 29 Strassmeier K.G. & Schordan P., 2000, AN 321, 277 Takeda Y., Sato B., Kambe E., Aoki W. et al. 2001, PASJ 53, 1211 Toner C.G., Gray D.F., 1988, ApJ 334, 1008 8 Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs Table 1 (Continued) HD HR 75732 76151 76780 81809 82106 86728 88072 89251 89269 89389 91347 95128 96094 98630 99491 101206 102870 107705 108954 109358 110833 110897 112758 114710 115383 116443 117043 117176 119802 122064 122120 124292 125184 126053 130322 131977 135204 135599 137107 139323 139341 140538 141004 143761 144287 144579 145675 146233 149661 151541 152391 154345 154931 157214 157881 158614 3522 3538 – 3750 – 3951 – – – 4051 – 4277 – – 4414 – 4540 4708 4767 4785 – 4845 – 4983 5011 – 5070 5072 – 5256 – – 5353 5384 – 5568 – – 5727 – – 5853 5868 5968 – – – 6060 6171 – – – – 6458 – 6516 Name 55 Rho1 Cnc 20 LMi 47 UMa 83 Leo 5 Bet Vir 17 Vir 8 Bet CVn 10 CVn 43 Bet Com 59 Vir 70 Vir 2 Eta CrB 23 Psi Ser 27 Lam Ser 15 Rho CrB 18 Sco 12 Oph 72 Her Teff this paper 5373 5776 5761 5782 4827 5735 5778 5886 5674 6031 5923 5887 5936 6060 5509 4649 6055 6040 6037 5897 5075 5925 5203 5954 6012 4976 5610 5611 4763 4937 4568 5535 5695 5728 5418 4683 5413 5257 6037 5204 5242 5675 5884 5865 5414 5294 5406 5799 5294 5368 5495 5503 5910 5784 4035 5641 N σ, K 97 88 87 85 76 91 82 89 95 48 75 89 73 52 96 60 48 56 60 72 80 68 83 71 40 83 98 104 71 84 35 89 89 79 85 62 91 86 60 90 90 100 81 81 93 89 98 96 90 88 82 87 82 85 9 98 9.7 3.0 5.0 6.9 6.0 5.6 5.0 6.3 5.7 8.9 7.4 3.8 11.6 10.0 8.6 7.6 6.8 7.8 5.5 6.2 3.9 12.3 8.4 6.8 9.3 9.9 4.7 4.7 6.6 8.1 11.4 4.0 5.9 6.9 5.4 6.8 4.6 5.1 6.9 7.7 7.9 3.5 4.4 11.1 5.7 10.3 12.1 3.8 3.2 6.4 4.5 5.6 6.7 9.5 4.5 3.6 Teff EDV93 Teff AAMR96 Teff BLG98 Teff DB98 5763 5611 5619 5746 5882 6176 6095 6060 5879 5867 4576 6124 6127 6068 6068 5959 5989 5862 5137 5985 5967 5795 6029 6021 5116 5964 5482 5562 4605 5937 5782 5635 5645 4609 4551 5897 5726 5309 5275 5676 4011 Mv B–V rem 5.456 4.838 5.011 2.945 6.709 4.518 4.717 3.292 5.089 4.034 4.725 4.299 3.725 3.043 5.230 6.750 3.407 4.104 4.507 4.637 6.130 4.765 5.931 4.438 3.921 6.175 4.851 3.683 6.881 6.479 7.148 5.311 3.898 5.032 5.668 6.909 5.398 5.976 4.237 5.909 5.115 5.045 4.072 4.209 5.450 5.873 5.319 4.770 5.817 5.630 5.512 5.494 3.558 4.588 8.118 4.910 0.851 0.632 0.648 0.606 1.000 0.633 0.593 0.569 0.645 0.532 0.513 0.576 0.550 0.553 0.785 0.983 0.516 0.498 0.518 0.549 0.938 0.510 0.791 0.546 0.548 0.850 0.729 0.678 1.099 1.038 1.176 0.721 0.699 0.600 0.764 1.091 0.742 0.804 0.507 0.943 0.898 0.640 0.558 0.560 0.739 0.707 0.864 0.614 0.817 0.757 0.732 0.708 0.578 0.572 1.371 0.678 * * * * * Kovtyukh et al.: Precise temperatures for 181 F–K dwarfs Table 1 (Continued) HD/BD HR 158633 159062 159222 159909 160346 161098 164922 165173 165401 165476 166620 168009 170512 171067 173701 176841 182488 183341 184385 184768 185144 186104 186379 186408 186427 187123 187897 189087 189340 190067 195005 197076 199960 201091 201092 202108 203235 204521 205702 206374 210667 211472 215065 215704 217014 219134 219396 220182 221354 +32 1561 +46 1635 Sun 6518 – 6538 – – – – – – – 6806 6847 – – – – 7368 – – – 7462 – – 7503 7504 – – – 7637 – – 7914 8041 8085 8086 – – – – – – – – – 8729 8832 – – – – – – Name 61 Sig Dra 16 Cyg A 16 Cyg B 11 Aqr 61 Cyg 61 Cyg 51 Peg Teff this paper 5290 5414 5834 5749 4983 5617 5392 5505 5877 5845 5035 5826 6078 5674 5423 5841 5435 5911 5552 5713 5271 5753 5941 5803 5752 5824 5887 5341 5816 5387 6075 5821 5878 4264 3808 5712 6071 5809 6020 5622 5461 5319 5726 5418 5778 4900 5733 5372 5295 4950 4273 5777 N σ, K 83 96 93 93 84 90 96 95 85 90 75 93 43 81 104 92 82 85 87 94 79 95 67 83 77 86 95 83 90 100 51 75 78 17 5 82 52 74 50 89 81 91 95 95 92 63 91 94 95 82 12 889 10.7 7.9 4.0 5.6 3.9 7.3 6.0 4.7 8.5 5.9 5.7 4.0 9.4 6.5 9.7 6.2 4.4 3.9 4.1 3.9 6.3 5.8 9.8 3.1 3.5 5.0 5.0 4.0 8.4 10.3 6.7 5.6 5.9 12.4 26.4 7.2 8.4 13.6 4.7 5.4 5.6 5.3 9.7 4.9 5.4 7.9 5.3 4.7 5.5 6.2 4.2 0.9 Teff EDV93 Teff AAMR96 Teff BLG98 9 Teff DB98 5770 5708 5852 4947 5781 4995 5833 4930 5826 5758 5227 5763 5767 5761 5813 4323 3865 5635 5755 4785 5783 5752 5774 5815 Mv B–V 5.896 5.485 4.653 4.459 6.382 5.294 5.293 5.388 4.880 4.406 6.165 4.528 3.965 5.191 5.343 4.487 5.413 4.201 5.354 4.593 5.871 4.621 3.586 4.258 4.512 4.433 4.521 5.873 3.920 5.731 4.302 4.829 4.089 7.506 8.228 5.186 3.606 5.245 3.839 5.304 5.470 5.835 5.131 5.500 4.529 6.494 3.918 5.661 5.610 6.493 7.895 4.790 0.737 0.706 0.617 0.657 0.950 0.632 0.789 0.732 0.557 0.580 0.871 0.596 0.542 0.660 0.847 0.637 0.788 0.575 0.721 0.645 0.765 0.631 0.512 0.614 0.622 0.619 0.585 0.782 0.532 0.707 0.498 0.589 0.590 1.158 1.308 0.610 0.468 0.545 0.513 0.674 0.800 0.802 0.594 0.795 0.615 1.009 0.654 0.788 0.830 0.919 1.367 0.65 rem * * *