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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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5839
5512
Mv
B–V
3.656
6.506
4.190
5.014
4.716
3.833
4.021
5.814
5.053
5.144
4.280
4.426
3.278
4.469
6.893
4.397
5.013
4.418
5.105
3.100
6.434
5.398
3.375
3.100
5.108
5.021
3.879
4.637
5.628
4.869
4.169
5.345
0.593
1.077
0.575
0.584
0.545
0.585
0.547
0.802
0.603
0.631
0.558
0.507
0.507
0.623
1.074
0.529
0.600
0.531
0.665
0.760
0.955
0.695
0.605
0.574
0.630
0.746
0.556
0.587
0.796
0.573
0.515
0.717
rem
*
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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
*
*
*