F.-J. Zickgraf 1 - J. Krautter 2 - S. Reffert 3 - J. M. Alcalá 4 - R. Mujica 5 - E. Covino4 - M. F. Sterzik 6
1 - Hamburger Sternwarte, Gojenbergsweg 112, 21029 Hamburg, Germany
2 - Landessternwarte Königstuhl, 69117 Heidelberg, Germany
3 - Sterrewacht Leiden, PO Box 9513, 2300 RA Leiden, The Netherlands
4 - Osservatorio Astronomico di Capodimonte, via Moiariello 16, 80131 Napoli, Italy
5 - Instituto Nacional de Astrofisica, Optica y Electronica, A. Postal 51 y 216 Z.P., 72000 Puebla, Mexico
6 - European Southern Observatory, Alonso de Cordova 3107, Santiago 19, Chile
Received 16 August 2004 / Accepted 3 December 2004
We present results of an investigation of the X-ray properties, age distribution, and kinematical characteristics of a high-galactic latitude sample of late-type field stars selected from the ROSAT All-Sky Survey (RASS). The sample comprises 254 RASS sources with optical counterparts of spectral types F to M distributed over six study areas located at , and . A detailed study was carried out for the subsample of 200 G, K, and M stars. Lithium abundances were determined for 179 G-M stars. Radial velocities were measured for most of the 141 G and K type stars of the sample. Combined with proper motions these data were used to study the age distribution and the kinematical properties of the sample. Based on the lithium abundances half of the G-K stars were found to be younger than the Hyades (660 Myr). About 25% are comparable in age to the Pleiades (100 Myr). A small subsample of 10 stars is younger than the Pleiades. They are therefore most likely pre-main sequence stars. Kinematically the PMS and Pleiades-type stars appear to form a group with space velocities close to the Castor moving group but clearly distinct from the Local Association.
Key words: surveys - X-rays: stars - stars: late-type - stars: pre-main sequence - stars: kinematics - solar neighbourhood
|Figure 1: Location of the six study areas in galactic coordinates. The dots show the positions of the RASS X-ray sources with stellar counterparts in the respective area. The solid and dashed curves denote the position and width, respectively, of the Gould Belt according to Guillout et al. (1998). In addition positions of several associations are shown.|
A first discussion of the properties of the late-type stellar component was given by Zickgraf et al. (1998) (Paper VI). A subsample of stars in study area I located some south of the Taurus-Auriga star forming region (SFR) was found to contain a large fraction of very young, presumably pre-main sequence stars. In order to investigate the age distribution of the complete sample of coronal X-ray emitters we obtained further low, medium and/or high resolution spectroscopic observations for most G-M stars in our sample. The goals were to carry out a lithium survey in order to identify lithium-rich high-galactic latitude G-M type stars and to determine precise radial velocities. In solar-like stars the lithium abundance can be used as an age estimator. Its knowledge therefore allows to study the age distribution of the X-ray active stellar sample. Combining the age information with proper motions and radial velocities would thus allow to investigate a possible age dependence of the kinematical properties of the stellar RASS sample.
This paper is structured as follows. The sample is presented in Sect. 2.
Observations and data reduction are described in Sect. 3.
Observational results are presented in Sect. 4.
Based on these results the sample properties are analysed and
discussed in Sect. 5.
Finally, conclusions are given in Sect. 6.
|Sep. 20-23, 1996||2100||CAFOS||CA 2.2 m|
|Jan. 31-Feb. 3, 1997||2100||CAFOS||CA 2.2 m|
|Jan. 2-6, 1998||1600||CAFOS||CA 2.2 m|
|Feb. 18, 1998||22 000||CASPEC||ESO 3.6 m|
|May 14-16, 1998||1300||DFOSC||ESO Danish 1.54 m|
|Dec. 22-25, 1998||34 000||FOCES||CA 2.2 m|
|Apr. 29-May 4, 1998||4600||CARELEC||OHP 1.93 m|
|Oct. 21-16, 1998||20 000||AURELIE||OHP 1.52 m|
|Jan. 11-15, 2000||34 000||FOCES||CA 2.2 m|
|Jun. 13-18, 2000||34 000||FOCES||CA 2.2 m|
|Dec. 3-6, 2001||34 000||FOCES||CA 2.2 m|
|Feb. 19-23, 2002||34 000||FOCES||CA 2.2 m|
For the spectroscopic follow-up investigation we selected the 200 X-ray sources from the catalogue with stellar counterparts of spectral types G to M. F stars were not included in the spectroscopic follow-up observations because for these stars the lithium abundance is not a good age estimator. In 19 cases two stars have been assigned as counterpart to the X-ray source in Paper III. Several of these secondary counterparts were also observed. As in Paper IV we will however only use the primary identifications for statistical purposes. The entire "coronal'' sample including the F type stars comprises 253 X-ray sources. The known RS CVn star HR 1099 (=V 711 Tau) which is X-ray source A031 in Paper III was excluded from the coronal sample discussed in the following. The sample finally selected for spectroscopic follow-up observations thus comprised 199 of the 200 X-ray sources with optical counterparts of spectral type G to M as listed in Paper III.
Infrared photometry in J, H, and K was taken from the Two Micron All Sky Survey
(2MASS) catalogue. From this data base infrared sources within 10
around the optical
position of the counterpart were extracted. A total of 267 2MASS sources was found of which 90% were located within 2
from the optical counterparts (including the 19 double
identifications, see above). We considered the 258 matches within 4
as reliable identifications. Matches between 4
individually checked and all found to be also correct. This means that for all but 5 RASS
sources (A035, A045, A065, D022, and D114) 2MASS measurements are available.
|Spec. type||This work||Paper IV|
The spectra were reduced with the standard routines of the ESO-MIDAS software package. The low- and medium-resolution spectra and the high-resolution spectra observed with AURELIE were reduced with the Longslit package. For the FOCES and CASPEC data the routines of the Echelle package were applied.
Spectra could be secured for the counterpart of 172 out of 199 RASS sources with spectral types between G and M. High resolution observations were obtained for 118 of the 141 G and K stars of the selected sample (originally 143 G-K stars minus A031 and E020). Lithium equivalent widths and radial velocities for six of the stars not observed by us with high resolution were adopted from high-resolution spectroscopic studies by Wichmann et al. (2001) (5 stars: A154, B049, B194, C062, C197) and Neuhäuser et al. (1995) (1 star: A058). Ten G-K stars fainter than 12th magnitude were observed only with low resolution. Thus for 134 of the 141 G-K stars spectroscopic follow-up observations exist. For the remaining 7 stars no observations could be obtained. Further high resolution data were found for the secondary counterpart of A098 in Favata et al. (1997). With a few exceptions M stars were observed with low resolution only. Due to bad weather conditions during the OHP observing campaign the M stars in area V could not be observed. In total 38 M stars were observed with low resolution and 7 with high resolution. For 13 M stars no observations could be obtained.
In the following we give more technical details of the spectroscopic observations.
A few stars were observed in May 1998 with the focal reducer camera DFOSC attached to the Danish 1.54 m telescope at ESO, La Silla. The spectra were obtained with grism No. 7 and a slit width of 1 . The wavelength range covered by the spectra was 3840-6845 Å. As detector the LORAL/LESSER CCD# C1W7 with a pixel size of 15 m was used. The resulting spectral resolving power was 1300.
In October 1998 high-resolution spectra were obtained with the spectrograph AURELIE at the 1.52 m telescope of the OHP. A description of the spectrograph can be found in Gillet et al. (1994). The spectra were observed with grating No. 2 with 1200 lines mm-1 giving a reciprocal linear dispersion of 8 Å mm-1. The detector was a double-barrette Thomson TH7832 (2048 pixel with 13 m pixel size). The spectra cover the wavelength interval from 6540 Å to 6740 Å. The resolution of the spectra is 20 000. Wavelength calibration was obtained with Neon and Argon lamps.
High-resolution spectra of 3 objects were obtained with the Cassegrain Echelle Spectrograph (CASPEC) at the ESO 3.6 m telescope on La Silla in February 1998. Wavelength calibration was obtained with a ThAr lamp. The CASPEC spectra cover the spectral range from 5350 to 7720 Å with a nominal resolving power of 22 000 (Sterzik et al. 1999).
During each high-resolution observing campaign radial and rotational velocity standard stars were observed in addition to the science targets.
During the observing runs a small set of spectroscopic standard stars, mainly of luminosity class V, had been observed together with the science targets. The coverage of the spectral type - luminosity class plane, however, was insufficient for a detailed two-dimensional classification. We therefore extended the spectroscopic data base for the standard stars by making use of the spectra available in the stellar library of Prugniel & Soubiran (2001) which is part of the HYPERCAT data base. We used the data set with a spectral resolution of 10 000. In order to match this resolution our FOCES, AURELIE, and CASPEC spectra were smoothed accordingly with an appropriate Gaussian filter. In this way the signal-to-noise ratio improved while the necessary spectral resolution for the classification was preserved. Spectral types and luminosity classes (LCs) of MK standard stars contained in the stellar library were adopted from Yamashita et al. (1976), Keenan & McNeil (1989), Garcia (1989), Keenan & Barnbaum (1999), and Gray et al. (2001). In a few cases we adopted the spectral classification given in Prugniel & Soubiran (2001). The grid of spectroscopic standard stars is listed in Table 3.
In a pilot study for the work presented here Ziegler (1993) studied the spectral
types of F, G and K-type stars from the RASS using spectra observed in the red spectral
6200-6750 Å). He found various line ratios useful for
classification purposes. For the F- and G-type stars the
ratios Fe I 6394/Si II
6346, Fe II6456/Ca I 6450 and
Fe II6456/ Fe I 6394 were found
to be good indicators for the spectral type. In K stars the ratios TiO
6240 / V I 6296
and Fe I 6250/Ca I 6450 were useful classification criteria.
|Star||Sp. type||Ref.||Star||Sp. type||Ref.|
|HD 222368||F7V||4||HD 188119||G7III||4|
|HD 016765||F7IV||6||HD 010700||G8V||4|
|HD 216385||F7IV||6||HD 188512||G8IV||4|
|HD 181214||F8III||6||HD 027348||G8III||3|
|HD 004614||G0V||4||HD 175306||G9III||4|
|HD 013974||G0V||4||HD 145675||K0V||5|
|HD 019373||G0V||4||HD 185144||K0V||4|
|HD 114710||G0V||1||HD 198149||K0IV||5|
|HD 150680||G0IV||4||HD 048433||K0III||3|
|HD 039833||G0III||6||HD 010476||K1V||5|
|HD 204867||G0Ib||1||HD 222404||K1IV||5|
|HD 204613||G1III||4||HD 096833||K1III||5|
|HD 185758||G1II||4||HD 022049||K2V||4|
|HD 186408||G2V||4||HD 137759||K2III||4|
|HD 126868||G2IV||2||HD 020468||K2II||4|
|HD 209750||G2Ib||1||HD 219134||K3V||4|
|HD 117176||G4V||5||HD 003712||K3III||3|
|HD 127243||G4IV||5||HD 201091||K5V||4|
|HD 186427||G5V||1||HD 118096||K5IV||6|
|HD 161797||G5IV||4||HD 029139||K5III||3|
|HD 027022||G5IIb||5||HD 088230||K6V||4|
|HD 206859||G5Ib||1||HD 201092||K7V||4|
|HD 003546||G6III||5||HD 079210||M0V||6|
|HD 182572||G7IV||4||HD 046784||M0III||6|
We used these ratios for the refinement of the spectral types given in Paper III. Figure 2 shows the histogram of the differences between the revised and original spectral types. The narrow peak shows that with few exceptions the overall agreement is good. We found a small mean difference of -0.5 subclass between the high- and low-resolution spectral types with a standard deviation of 2.2 subclasses. The original and the revised statistics of spectral types are listed in Table 2. In nine cases the difference of the spectral types was larger than 3 subclasses. The largest differences were found for B174 and E256 (-6 subclasses), B185 (7 subclasses), D018 (9 subclasses), and E022 and E067 (-9 subclasses). The LFOSC spectrum of E256 was actually classified as K4, but erroneously entered in Paper III as M0. For D018 which is a very bright star the original LFOSC spectrum classified as G2V could suffer from saturation. In SIMBAD this star is listed as K0III (Schild 1973). The classification based on the FOCES spectrum is K1III, which is in good agreement with the literature. We adopt this spectral class in the following. For the remaining stars with large deviations no LFOSC classification spectra were obtained. The spectral classes were adopted from SIMBAD. In the following we use the improved FOCES classifications.
Following Gahm & Hultqvist (1972) and Ziegler (1993)
luminosity classes (LC) were obtained using the strength of the lines of
5854 Å, 6497 Å, Sc II 6605 Å, and
La II 6390 Å. We added the Y II 6614 Å line which also shows a clear
The ratio of Sc II 6605 Å and Y II 6614 Å is a good
luminosity indicator for spectral types earlier than about K5-7. For spectral types
later than K0 the
strength of La II was additionally useful to discriminate luminosity classes III and
higher from LC V and IV. For G stars LC III and higher could also be discriminated
from LC IV by the use of this line.
Comparing in this way the line strengths and ratios in the MK standards
with the sample stars LCs could be assigned to most stars. For a few stars the stellar
absorption lines were strongly broadened by rapid rotation (see below). In these cases it
was not possible to determine the luminosity class due to the
limited S/N of the spectra and to line blending. The limit was reached around
km s-1. For the rapid rotators we adopted LC V. As discussed in
Sect. 5.1.1 we used the luminosity classes to derive spectroscopic
|Figure 2: Comparison of the spectral types derived from the classification spectra used in Paper III (Sp(old)) and from the new high resolution spectra (Sp(rev.)). The abscissa is the difference (in spectral classes) between the revised and the original spectral types.|
The width of the cross-correlation function is a measure for the rotational
We therefore calculated the cross-correlation function as
before but for rotational velocity standards.
Standard stars with low
and spectral type as
close as possible to that of the objects were used for the cross-correlation analysis
as well as to calibrate the FWHM vs.
relation. From the FWHM of the
was then determined following the method described
in Covino et al. (1997).
Observations of rotational standard stars yielded a detection limit of
about 5 km s-1. From the statistics of the differences between
measured rotational velocities of rotational standard stars and
literature an uncertainty of
of 3 km s-1 could be estimated.
For rotational velocities above 40 km s-1 the shape of the peak of the
correlation function deviates increasingly from a Gaussian leading to larger
errors of 5-10 km s-1.
Figure 3 shows the histogram of the rotational velocities which
are listed in Table A.2.
In the high-resolution spectra the equivalent widths were measured directly by integrating the flux in the normalized spectra. The contribution of the neutral iron line Fe I Å was corrected according to the procedure described by Soderblom et al. (1993b). For stars with rotational velocities larger than 30 km s-1 the contribution of the Fe I lines near Li I was corrected in the following way. From the stellar library of Prugniel & Soubiran a spectroscopic standard star with a spectral type as close as possible to the target was selected. It was folded with the appropriate rotational velocity to match the broadened lines of the target spectrum. Then the EW of the Fe I absorption features was measured in the same wavelength interval as used to determine the Li I EW in the target spectrum. Finally the corrected lithium EW was obtained by subtracting the contribution of the Fe I lines from the measured lithium EW of the target spectrum. Errors of the high-resolution EWs are typically 5-15 mÅ, depending on the signal-to-noise ratio and on the rotational velocity. The EWs are listed in Table A.4.
In Fig. 4 the EWs obtained from the low- and
the high-resolution spectra are compared. In the low-resolution spectra the
EWs W(Li I ) are obviously slightly underestimated by about 40 mÅ. However,
the overall agreement is good
and the differences are only of the order the uncertainty of the low-resolution measurements.
This demonstrates that the fitting method applied to the low-resolution spectra works
remarkably well. In particular, W(Li I ) is not overestimated as it would be the case
if the EWs would be determined directly by flux integration without taking the contribution
of the Fe I lines into account.
|Figure 4: Comparison of the equivalent widths of Li I determined from the low- and the high-resolution spectra. The dashed line denotes a ratio of 1 of the two measurements.|
In a few binaries lithium lines could be identified in one or both components. In order to disentangle the lines of the individual components and to identify a possible Li I line spectra from the Prugniel & Soubiran sample with the appropriate spectral types were folded with the rotational profile for the measured and shifted with respect to the measured radial velocities. Then the spectra were superimposed by using appropriate values for the relative flux contributions. Finally the resulting artificial binary spectrum was compared with the observed spectrum. Correction factors for the measured lithium equivalent widths were estimated from the artifical spectrum. In most cases the spectra suggest a flux ratio of 1 to 2 for the individual components at 6708 Å. Exceptions are e.g. A001 and A071. In A001 the primary component is a fast rotator ( km s-1) whose broad lines dominate the spectrum. Of the secondary component only the strongest lines of a mid to late type K star are detectable. For this binary system we adopted a flux ratio of 5:1 for the continuum contributions of the primary and secondary component at 6708 Å. In A071 both components are fast rotators with very broad lines. In this case it was not possible to determine a lithium EW for each component. The total EW was therefore assigned in equal shares to the individual components and the lithium equivalent widths were corrected by assuming equal flux contributions. The triple system B160 is even more complicated. It consists of 3 early to mid G-type stars with spectral types between G2 and G5. Two of the three components exhibit a lithium absorption line.
It is clear that the equivalent widths of the binaries and the triple system are less reliable than those of the single stars due to the uncertainty of the continuum correction. In Table A.4 the lithium EW of the strongest component is given.
For 74 stars a spectroscopic parallax could be derived from the high-resolution spectra
by adopting the absolute V magnitudes, as appropriate for the spectroscopically
determined luminosity class, from Schmidt-Kaler (1982).
For the bulk of M stars we used infrared JHK measurements from the 2MASS catalogue
to derive a photometric distance. The two-colour diagram of J-H and H-K is displayed
in Fig. 5. It shows that the M stars
are distributed around the locus of main-sequence stars (solid line in
Fig. 5). For the further analysis distances of M stars
were therefore estimated by adopting MV for LC V from Schmidt-Kaler
(except for the 11 stars with trigonometric parallaxes).
This adds 43 more RASS sources with a distance estimate.
Thus total distances are available for 100 G-K and 54 M stars.
For the remaining
stars without a distance measurement we derived a lower limit for the
distance by assuming that they are main-sequence objects with LC V.
|Figure 5: Two-colour diagram for the infrared magnitudes from 2MASS. Circles denote M stars, crosses stars with spectral types F to K. The solid, dotted, and dashed lines denote the loci of main sequence stars, giants, and supergiants, respectively.|
An estimate of the error of the spectroscopic and photometric distances, , may be obtained from the following considerations. The error is due to the uncertainties of the absolute visual magnitude, MV, and of V. For the latter we conservatively adopted the error of the photographic GSC magnitudes for all stars. The dominating source of uncertainty is the error of MV. For G-K stars of LC V and IV and correspondingly for LC III and II we used half of the difference of MV of these luminosity classes as estimate for . This leads to an estimate for of 30-50%. In the case of M stars the main source of error of MV is due to the uncertainty of the spectral class. This also leads in total to % if an uncertainty of 1-2 spectral subclasses is assumed. We finally adopted 50% as relative error for spectroscopic and photometric distances.
For the derivation of the distances interstellar extinction was not taken into account. Given the high galactic latitude of our sample it is actually expected to be small. With the relation given by Spitzer (1978) with the column density of neutral hydrogen, , and colour excess E(B-V) upper limits of the extinction can be estimated. We expect extinction values, , of less than 0.2-0.3 in all study areas except area I. This region could have a higher extinction of up to 0.6 magnitudes for the most distant stars. For these estimates the values given in Paper II were used.
For 20 stars in our sample both spectroscopic and Hipparcos parallaxes,
exist. They are compared in Fig. 6.
The agreement of the two distance measurements
for this subsample is good. The mean ratio of both parallaxes is
For the further analysis we adopted the spectroscopic parallaxes if no Hipparcos
or other trigonometric
parallax was available. The adopted distances are listed in
|Figure 6: Comparison of the distances derived from spectroscopic or photometric parallaxes and from trigonometric parallaxes, . The dashed line denotes equal distance values. M stars with distance estimates from the 2MASS IR photometry are plotted as circles.|
Figure 7 shows the number distribution of the distances for the 184 F-, G-, and M stars. Also shown is the distribution including the stars with minimum distances estimated by adopting LC V. The number distribution of the total sample has a maximum around 50 pc with a tail extending up to several 100 pc. Most stars are nearer than 200 pc, 33 stars have distances above 300 pc (including 16 stars with minimum distances), and in 4 cases (not shown in Fig. 7) we derived a distance above 1 kpc (including 3 stars with minimum distances). The identifications of the very distant RASS counterparts may be questionable.
For the stars with trigonometric parallaxes the absolute magnitude, MV, was calculated from the distance and visual magnitude given in Table A.2. A luminosity class was then assigned according to Schmidt-Kaler (1982). Likewise, bolometric corrections were taken from the same reference to determine the bolometric magnitudes for all stars with known distances.
As expected the majority of stars with a luminosity class determination, 90%,
have luminosity class V or IV. A small number of 17 stars was classified as giants
(LC III-IV, III, and II), 12 of these based on Hipparcos parallaxes.
In Fig. 8 the H-R diagram is shown for
all stars with a spectroscopic or trigonometric parallax.
M stars are shown only if a trigonometric parallax
|Figure 7: Histogram of the distance distribution. The solid lines represent the distribution of trigonometric, spectroscopic, and photometric parallaxes. The dashed lines include distance estimates derived from assuming absolute visual magnitude of main-sequence stars for the remaining stars without other distance estimate.|
|Figure 8: H-R-diagram for single stars with either a trigonometric parallax from Hipparcos or other sources (+ sign) or with a spectroscopic parallax (triangles).|
|Figure 9: X-ray luminosity for single stars vs. distance. + signs mark stars in study areas I, II, III, IV, and VI, signs represent stars in area V. The solid and dotted lines mark the flux limits for the two groups of study areas.|
|Figure 10: X-ray luminosity for single stars as a function of effective temperature, . Different symbols identify stars with trigonometric (+ sign), spectroscopic (triangles), or IR photometric parallaxes (circles). The variability range of solar X-ray emission in the ROSAT-PSPC pass band is marked by the vertical bar.|
|Figure 11: X-ray luminosity for all single stars with trigonometric (+ sign), spectroscopic (triangles), or IR photometric parallaxes (circles) as a function of absolute visual magnitude, MV. The variability range of solar X-ray emission is marked by the vertical bar.|
In Fig. 10 is plotted versus the effective temperature and Fig. 11 shows as function of the absolute visual magnitude, MV. A weak trend of increasing with increasing is visible. The -MV diagram shows a clear correlation with decreasing for decreasing optical luminosity. This reflects the fact that depends on the emitting surface. The width of the -MV distribution at a given MV tells that the X-ray surface flux density of the stars in our sample spans a range of a factor of 1000. Around MV = 5 the lower limit of the X-ray luminosities of the sample stars is about a factor of 10 above the solar soft X-ray variability range ( erg s-1, Schmitt 1997). The upper limit of in our sample is about a factor of 10-30 higher than in the volume-limited sample of Schmitt (1997).
The ratio of and bolometric luminosity, , is plotted in Fig. 12 as function of . A clear correlation is visible with the low luminosity stars with later spectral types having the highest ratio of . This is in agreement with the results of Fleming et al. (1995) who studied the coronal X-ray activity of low-mass stars in a volume limited sample. They found the highest ratios of for dMe stars. As discussed in Paper IV, most M stars in our sample are actually dMe stars, that is of the 58 M stars listed originally in Paper III 53 exhibit H emission lines. Note, however, that selection effects inherent in our flux-limited sample may also play a role.
The X-ray surface flux density is displayed as a function of MV in Fig. 13 and
as a function of
in Fig. 14. Our sample contains
mainly stars with a high surface flux density which is on the average 1 to 2 orders of
magnitude above the solar flux level. This can be understood in view of the result
discussed below in Sect. 5.2.2 that our sample contains a large
fraction of young and hence very X-ray active stars.
Old solar-like stars are obviously not present in our sample. The maximum value
of the surface flux density of our sample stars is around
108 erg s-1 cm-2. This value is consistent with the result obtained by
Schmitt (1997) who found a maximum around
107-108 erg s-1 cm-2 in his volume-limited sample of solar-like stars.
|Figure 12: Ratio of X-ray and bolometric luminosity for all single stars with trigonometric (+ sign), spectroscopic (triangles), or IR photometric parallaxes (circles) as a function of bolometric magnitude, .|
|Figure 13: X-ray surface flux density for all single stars with trigonometric (+ sign), spectroscopic (triangles), or IR photometric parallaxes (circles) as a function of absolute visual magnitude, MV. The vertical bar marks the typical flux level of solar coronal holes in the ROSAT-PSPC pass band.|
|Figure 14: X-ray surface flux density for all single stars as a function of effective temperature, . The meaning of the symbols is the same as in Fig. 10. The vertical bar marks the typical flux level of solar coronal holes in the ROSAT-PSPC pass band.|
Finally, in Fig. 15 the ratio
is displayed as
projected rotational velocity, .
No clear correlation can be seen, except that small ratios of
are only found
for small ,
whereas fast rotators exhibit high
The lithium equivalent widths were converted to abundances, N(Li),
by using the curves of
growth of Soderblom et al. (1993b) for stars with
and of Pavlenko & Magazzù (1996) and Pavlenko et al.
(1995) for cooler stars. As in Paper VI effective temperatures
were derived from the spectral types using the temperature calibrations of
de Jager & Nieuwenhuijzen (1987).
The uncertainty of
is typically 200 K. This leads to errors of the estimated Li
abundances of about 0.3 dex. Lithium abundances are shown in
Fig. 17 as function of effective temperature with
indicated by the symbol size.
|Figure 16: Equivalent widths of Li I as a function of for all stars in the six study areas. Crosses and triangles are high and low resolution measurements, respectively. The solid lines represent the upper and lower envelope of the lithium equivalent widths in the Pleiades adopted from Soderblom et al. (1993b). The dashed line shows the upper envelope for the Hyades cluster taken from Thorburn et al. (1993).|
|Figure 17: Lithium abundances versus effective temperature for the complete sample. Upper limits are plotted as downward arrows. Circles denote high-resolution measurements with the symbol size depending on . Low and medium resolution data are plotted as triangles. The solid lines are the upper and lower limit of (Li I ) in the Pleiades; the long dashed and short dashed lines show the upper (Li I ) limits for the UMaG and the Hyades, respectively.|
In Fig. 16 the upper and lower envelopes of the (Li I ) distributions for stars in the Pleiades and the upper envelope for the Hyades are shown. Likewise, Fig. 17 includes the upper envelopes of the lithium abundances of stars in the Pleiades, the UMaG, and the Hyades, and in addition the lower envelope for the Pleiades.
Using the lithium abundance data for the mentioned clusters and moving groups we finally defined four age groups. The age group "PMS'' consists of stars above the Pleiades upper envelope and is thus younger than the Pleiades, i.e. younger than 100 Myr. The group of stars between the upper and lower Pleiades envelopes can be assumed to have an age similar to the Pleiades. In the Pleiades the G and K stars are supposed to have reached the ZAMS. This group with an age of 100 Myr is therefore designated "Pl_ZAMS''. The age group "UMa'' comprises stars between the lower Pleiades and the upper Hyades envelope. The age of the stars of this group is between 100 and 600 Myr, i.e. on the average 300 Myr, which is the age of the UMaG. The age group "Hya+'' comprises G-K stars with either a lithium abundance below the upper Hyades envelope or with an upper limit for the lithium abundance only. The latter means that this group also contains stars for which the upper limit is above the Hyades line. Evolved stars more luminous than LC IV are included in the age group "Hya+'' if not stated otherwise in the following. It should be noted, however, that due to the well-known scatter of the lithium abundances in clusters stars below the upper envelope for the corresponding age group are not necessarily older than the respective group. Therefore, the "Hya+'' group might actually also contain some younger stars although it certainly is dominated by truly old stars.
In M stars older than several 106 yr lithium has been destroyed already (e.g. D'Antona & Mazzitelli 1994). With the exception of two stars we could not detect lithium in the M stars of our sample. This means that the M stars are typically older than 10 Myr. We thus only defined a group "M stars'' without assigning an age. This group does not contain the two lithium rich M stars (see below). We will return to the M stars in Sect. 5.3.1 where we use the kinematical properties to estimate their age.
Figure 17 shows that a small but significant group of 12 stars exists above the Pleiades upper limit. These objects appear thus to be younger than 100 Myr and may be even younger than or comparable to the age of IC 2602, i.e. 30 Myr. Two of these stars, B002 and F0140, are however giants (LC III) and are therefore not pre-main sequence (PMS) but evolved objects. This leaves a group of 10 stars which appears to consist of PMS objects, i.e. true members of the age group "PMS''. Actually, 8 of these 10 stars are found in area I which is located south of the Tau-Aur SFR. They represent the young stellar population in this region discussed in Paper VI. The remaining two stars are located in area II. The subsample of the lithium-rich stars including the giants is listed in Table 5. Their high-resolution spectra are shown in Fig. 18 except for A058. The spectrum of this star can be found in Neuhäuser et al. (1995). For its low-resolution spectrum see Paper VI. The spectrum of the M4 star B026 is displayed separately in Fig. 19.
The rotational velocities of the Li-rich stars are high on the average. Only the
below 10 km s-1. Six of the ten PMS stars have
km s-1. Table 4 lists the median
for each age
group. It shows that
decreases on the average with increasing age.
|Figure 18: Spectra of the lithium-rich sample listed in Table 5. The wavelengths of Li I Å and Ca I Å are indicated by the dashed lines.|
|Figure 19: Low-resolution spectrum of the M4 star B026. The dashed lines indicate Li I Å and Ca I Å.|
|field||RASS name||Sp. type||EW(Li I )||(Li)|
|field||RASS name||Sp. type||EW(Li I )||(Li)|
|A126||RX J0405.6+0341||G0V-IV||63||2.54||< 5|
The majority of stars has EWs and lithium abundances below the Pleiades upper limits of EW and (Li), respectively. In the region between the upper and lower envelope of the Pleiades 43 G-K stars are found. This group is listed in Table 6. Three of these stars are giants with LC IV-III, III, and II. The 40 non-giants appear to constitute a population with an age similar to the Pleiades, i.e. 100 Myr. The region between the Hyades upper and the Pleiades lower envelope contains 23 stars of which 4 are evolved objects. The UMa age group with an age of 300 Myr thus consists of 19 stars. Below the upper limit of the Hyades 57 non-giant stars are found and are thus assigned an age of older than 600-700 Myr. Adding the 17 evolved G-K stars which are certainly also older than 1 Gyr results in a total of 74 stars for age group "Hya+''. Thus lithium abundances and luminosity classification suggest that 47% of all G-K stars in the sample have an age of less than about 600-700 Myr. Restricting these statistical considerations to the later spectral types increases the fraction of stars younger than the Hyades. Of the 114 G5-K9 stars 55, i.e. 50%, have a lithium abundance higher than the Hyades. With the above mentioned ambiguity of the age group definition this means that at least half of the G5-K9 stars are younger than the Hyades. Some statistics of the age distribution of our sample stars for these age groups is summarized in Table 7.
As expected area IV located near the north galactic pole has the lowest surface density of stars younger than the Hyades. In this area only 2 stars younger than 600-700 Myr are found in 72 deg2. This corresponds to a surface density of deg-2 at a RASS count-rate limit of 0.03 cts s-1. In the other 5 areas (613.2 deg2) a total of 60 stars (including 5 stars in area V above 0.03 cts s-1) yields a surface density of deg-2. Counting stars of all age groups area IV has a surface density of deg-2 compared to deg-2in the other areas at the same count-rate limit. A t-test shows that these differences are significant.
The very young stars of the PMS sample are apparently more abundant in area I
than in any other area: 80% of these stars are found in area I.
Adding up the numbers of stars younger than the Hyades in areas II, III, and VI
leads to an average surface density
This is less than half
of the value in area I which is
Although indicative for a higher concentration of young stars in
area I the difference is not significant.
|<100 Myr||100 Myr||300 Myr||>660 Myr|
|<100 Myr||100 Myr||300 Myr||>660 Myr|
We compared the observed cumulative number distribution, , of our sample with model predictions by Guillout et al. (1996). The median latitude for the combined areas I, II, III, V, and VI, which are distributed between galactic latitudes of 20 and 50 , is actually , thus matching this model parameter well. The models of Guillout et al. (1996) give cumulative surface densities, , as a function of ROSAT-PSPC count rate, S, for three age bins: age younger then 150 Myr, age between 150 Myr and 1 Gyr, and older than 1 Gyr. We restricted the comparison to the youngest model age bin and to the sum of all model age bins because of the difficulty to separate observationally stars with ages of several 100 Myr to 1 Gyr and older. We further considered the combined sample of G and K stars. M stars were not included because of the lack of an observational age determination for stars of this spectral type in our sample. The uncertainty of the ages derived observationally from lithium was taken into account by forming two observational age samples matching as closely as possible the youngest age bin of the models: a) a sample comprising the sum of G-K stars from the PMS and Pl_ZAMS age group, and b) a sample containing in addition the corresponding UMa stars. The true sample of stars younger than 150 Myr is expected to lie between these limits.
The result of the comparison of
is depicted in
Fig. 20 for three RASS X-ray count rates of 0.1, 0.3 and 0.01 cts s-1. The predicted numbers of G-K are in good agreement
with our sample in the 5 study areas located around
in galactic latitude. This holds for both the sum of all age groups and stars younger than
150 Myr obtained as described above and represented in the figure
by the filled symbols. Likewise, the predicted flattening of
at lower count rates is also found in our data for area V which has the lowest
count rate limit of 0.01 cts s-1.
|Figure 20: Comparison of observed co-added number densities of G and K stars, , for three RASS count rates S with models of Guillout et al. (1996) for . Open symbols denote the sum of all age groups. Lower and upper filled symbols represent the sum of age groups "PMS'' and "Pl_ZAMS'', and of "PMS'', "Pl_ZAMS'', and UMa, respectively.|
|Figure 21: Proper motions for the six study areas. The different symbols denote the different age groups: filled circles = PMS, filled triangles = Pl_ZAMS, open circles = UMa, open triangles = Hya+, = giants, + = M stars without Li detection.|
Altogether, we were able to assign proper motions to the counterparts of 129 RASS sources with spectral types G to M. In detail we found 55 of 56 G stars, 61 of 86 K stars and 13 of 56 M stars in the mentioned catalogs. In addition we also found 54 F stars. An equal number of proper motions comes from Tycho-2 and UCAC2, while only two proper motions each were taken from Hipparcos and TRC, and only one each from PPM and STARNET, while the ACT was not used in the end at all.
These proper motion data were supplemented for the optically faint stars (mainly of spectral type K and M) by data from other catalogs: 53 stars from USNO-B1.0 (Monet et al. 2003, 36 M stars, 16 K stars, and 1 G star), 1 M star from Carlsberg Meridian Catalogs (1999), and 1 M star from the NPM1 Catalog of the Lick Northern Proper Motion Program (Klemola et al. 1987). Note that the USNO-B1.0 proper motions are not absolute, but relative to the Yellow Sky Catalog YS4.0 in the sense that the mean motion of objects common to USNO-B1.0 and YS4.0 was set to zero in USNO-B1.0. According to Monet et al. (2003) the difference between these relative proper motions and the true absolute ones should, however, be small.
Thus in total proper motion data are available for all G stars, for 77 of 86 K stars, and for 51 of 56 M stars.
The proper motions are shown in Fig. 21 for the individual study areas. The diagram displays proper motions for six groups of stars, i.e. the "PMS'' "Pl_ZAMS'', "UMa'' and "Hya+'' age groups, evolved stars (giants) and the M stars without lithium detection.
Of particular interest are the proper motions of the stars of the youngest age groups in area I, i.e. the "Pl_ZAMS'' and "PMS'' samples with ages of 100 Myr and less, This study area is located near the Tau-Aur SFR and near the Gould Belt (see Fig. 1) and has, probably due to its location, the highest surface density of young stars. Proper motions exist for all eight young stars of area I listed in Table 5. They are plotted in the upper left panel of Fig. 21. Four of the stars of age group PMS in area I were already identified as pre-main sequence objects by Neuhäuser et al. (1995) (A058, A069, A090, and A104). They assigned an age of 35 Myr to these stars. Likewise, one star of the Pl_ZAMS sample, A120, was assigned an age of 100 Myr by Neuhäuser et al. These stars were part of a sample investigated kinematically for membership to the Taurus-Auriga SFR by Frink et al. (1997). They studied stars in the central region of Tau-Aur and in a region south of Tau-Aur which partially overlaps at the southern edge with our area I. The sample studied by these authors contains three further stars of our sample, A007, A107, and A122, for which Neuhäuser et al. assigned an age of older than 100 Myr. This is in agreement with our age estimate of older than 660 Myr for A007 and A107, and of 300 Myr for A122.
In Table 9 the mean proper motions and their dispersions are summarized for the PMS, Pl_ZAMS, UMa, and Hya+ age groups, and for M stars without Li I detection. Obviously, the 8 "PMS'' stars show a smaller spread in proper motions than the older stars. They cluster around ( , ) of (+16, -8) mas yr-1 with a scatter of 15 mas yr-1 in each direction. Frink (1999) transformed the proper motions given by Frink et al. (1997) from the FK5 to the Hipparcos system and determined mean values of ( +8.7,-11.2) mas yr-1 for the southern sample of Frink et al. (1997). For the central region of Tau-Aur Frink (1999) derived mean proper motions of ( +4.5,-19.7) mas yr-1. The comparison of our results with the findings of Frink (1999) reveals an interesting trend in the mean proper motions relative to the core region of Tau-Aur. The southern sample of Frink et al. moves away from the centre of Tau-Aur with a mean proper motion of (+4.2, +8.5) mas yr-1. The PMS stars in area I are located even more to the south of the centre and their relative mean proper motion is actually even larger, (+12, +12) mas yr-1. Thus we find that the stars in area I move in approximately the same direction as the southern stars of Frink et al., but with an even higher proper motion.
Inspection of Fig. 2 in Frink et al. (1997) allows to
estimate a dispersion of about 15 to 20 mas yr-1 for both subsamples
which again is compatible with the 15 mas yr-1 derived for our PMS subsample.
The Pl_ZAMS stars exhibit a dispersion of the
proper motion which is larger
by a factor of 2 to 3. On the other hand, the UMa sample though being older
shows more coherent proper motions with a dispersion equal to the PMS stars.
The old stars of the Hya+ group and the M stars exhibit the largest dispersions.
Similar results are found for the other study areas.
So far we have considered the proper motions which depend on the distance and contain a contribution due to the solar motion. We therefore calculated tangential velocity components, vl and vb, in galactic coordinates, l and b, by using the distance estimates discussed above and the relations km s-1 and km s-1, with and being proper motions in galactic coordinates given in arcsec yr-1 and the distance d in pc. A table summarizing the resulting velocities and their dispersions for the individual study areas can be found in the Appendix (Table A.1).
The direction-dependent part of the tangential velocities due to the solar reflex motion
can finally be removed by transforming these velocities to the local standard
of rest (LSR). This is achieved by adding the corresponding solar velocity components.
We used the solar motion vector of Dehnen & Binney
) = (+10.0, +5.25, +7.17) km s-1 (see
below for the definition of the space velocities) to determine the solar reflex motion:
In Table 10 mean proper motions in galactic coordinates with respect
to the LSR,
the corresponding tangential velocities,
are listed for the different age groups. As discussed before the
exhibit the largest dispersion of the proper motions. Taking the distance effect into account the
dispersions of the respective tangential velocities are reduced to values similar to those obtained
for the Pl_ZAMS and UMa age groups. This again leads to the conclusion that the M stars have
on the average an age of 100-600 Myr. The largest velocity dispersions are found
for the Hya+ age group.
|Figure 22: Space velocities U, V, and W in the LSR frame. Stars of age groups "PMS'' and "Pl_ZAMS'' are plotted as filled circles. Open circles denote stars of the UMa and "Hya+'' age group. Giants are plotted as asterisks.|
|Figure 23: Upper panel: U-V velocity diagram for the youngest age groups "PMS'' (circles) and "Pl_ZAMS'' (triangles). The solid line encircles the region defined by Eggen (1984, 1989) to contain the young disk population. Also shown as large crossed circles are the U and V velocities of the Hyades supercluster, the Local Association (designated "local''), the Castor MG, and the UMa MG. Lower panel: W-V diagram for the same sample of stars. All velocities are in the LSR reference frame.|
The space velocities components are plotted in Fig. 22. The plot contains stars of all age groups and also includes the evolved stars (giants). Figure 23 shows in an enlarged scale the V-U, and V-W diagrams for the two youngest stellar age groups only, i.e. PMS and Pl_ZAMS stars.
As can be seen in Fig. 22 the filled symbols representing the youngest age groups, PMS and Pl_ZAMS, are more concentrated than the open symbols and the asterisks denoting the older age groups and giants, respectively. This can be tested by various statistical methods. First, we combined on one hand the PMS and Pl_ZAMS samples and on the other hand the older stars and giants in order to create distributions of the space velocity for the young and the old stars, respectively. A one-dimensional two-sample Kolmogorov-Smirnov (K-S) test on these distributions yields a probability of < that they are drawn from the same parent distribution. Likewise, the K-S test on the PMS and the complementary non-PMS sample yields a probability of only for having the same distribution. Therefore, PMS and non-PMS stars also have different space velocity distributions. Contrary to this, with a probability of 0.31 PMS and Pl_ZAMS stars have the same distribution. An F-test on the individual velocity components U, V, and W of the combined PMS-Pl_ZAMS and the older age groups shows that with a very low probability P their distributions are drawn from the same parent distribution, namely PU = 0.006, PV= 0.02, and . In particular, the velocity component perpendicular to the galactic plane, W, is significantly different in the young and the old age groups (see below).
In the following we will discuss mean velocities and velocity dispersions of the
different age groups.
These were calculated
as maximum-likelihood (M-L)
which takes into account that the measurement errors are
different for each star.
Following Pryor & Meylan (1993) M-L estimates of the mean
of U, V and W
were obtained together with errors by assuming that the velocities are drawn from
a normal distribution
In Table 11 the mean space velocities and velocity dispersions of the different age groups are summarized. Clearly the PMS sample has the smallest velocity dispersions. For stars with weak or no lithium detection the dispersions are the largest. The "Hya+'' subsample contains a significant fraction of older disk stars. This is particularly evident for the velocity component perpendicular to the galactic plane, W. Its dispersion increases from 2 km s-1 for the PMS sample to 30 km s-1 for the old lithium weak sample. The increasing velocity dispersion with increasing age reflects the effect of disk heating in the galaxy.
The PMS subsample in particular exhibits M-L mean space velocity components km s-1 and velocity dispersions of ( ) km s-1. This suggests that the PMS stars are kinematically related and may even form a kinematical group, but of course, the sample is small and the indicated relation should be considered more as a working hypothesis to be tested with extended samples. At this point it should be noted that the PMS star B206 in area II interestingly has space velocity components similar to the stars in area I. Unfortunately, no high resolution RV measurement is available for the second PMS star in area II, the M4Ve dwarf B026. In order to obtain at least an estimate of its space velocity components we measured the radial velocity using the low-resolution CAFOS spectra and the emission lines of H, H, H, and Ca II K. This yielded km s-1 with an error of about 20 km s-1. The resulting space velocity components are km s-1, km s-1, and km s-1. Within the errors the velocities of B026 are consistent with the mean velocities of the PMS sample. But clearly, a more accurate RV measurement is needed for B026 to confirm that both Li-rich stars in area II belong to the same kinematical group as the corresponding stars in area I as indicated by the presently available data. Note also from Fig. A.1 that in areas I and II the numbers of Li-rich stars are higher than in the other areas.
Enlarged sections of the U-V and W-V diagram are shown in Fig. 23 for the 35 stars of the two youngest age groups with measured
space velocities. The U-V diagram in the upper panel
includes the limits of the region occupied by the young disk stars as defined by Eggen
(1984, 1989). Indeed, as expected for a young stellar sample many,
albeit not all, stars have (U,V) velocities inside Eggen's box.
Also indicated are the (U,V) velocities of several
young stellar kinematical groups:
the Hyades supercluster, the Ursa Major moving group (UMa MG), the Local
Association (Pleiades MG), and the Castor moving group (Castor MG)
(for references see e.g. Montes et al. 2001).
Figure 23 suggests the existence of a kinematical subgroup in the combined PMS and Pl_ZAMS sample, which contains 35 stars with measured U, V, and W velocities. The group is concentrated near the velocity of the Castor MG at the upper V limit of Eggen's disk stars with 17 of the 35 stars found within a radius of 10 km s-1 around the velocity of the Castor MG. Six of these belong to the age group PMS and the rest to the Pl_ZAMS group. The M-L mean velocities of the subgroup are km s-1, and the velocity dispersions are km s-1. In V the group of 17 stars is somewhat off the Castor MG for which Palous & Piskunov (1985) give ( km s-1. Given the relatively small number of data points we may ask whether this concentration is due to a chance coincidence in an actually random distribution. We tested this possibility for the null hypothesis that the true underlying distribution of velocities in the U-V plane is random within a given circle around the origin. In a Monte Carlo simulation we calculated a large number of random velocity vectors in the U-V plane and counted the number of cases in which we found 17 stars within 10 km s-1 around the Castor MG velocity. For a random velocity distribution within a radius of 35 km s-1 containing 90% of the 35 stars, i.e. 31 stars, these simulations showed that we can reject the null hypothesis on a high significance level of >99.8%. The test radius of 35 km s-1 may be too small because it excludes 10% of the stars. Increasing the radius leads however to even higher significance levels. Decreasing the radius only leads to significance levels of <99% if the random distribution is calculated for radii smaller than 20 km s-1 which contains 65% of the PMS-Pl_ZAMS stars. Therefore we are lead to the conclusion that with a very high probability the concentration of velocities vectors in the U-V plane is not a chance coincidence.
An interesting feature is the accumulation of the 6 "PMS'' stars around a mean velocity of ( km s-1. This is not far from the velocity of the Castor MG (see above), but clearly distinct from the Local Association which has ( U,V) = (-1.6, -15.8) km s-1 (Montes et al. 2001). The velocity dispersions of this subgroup of PMS stars are km s-1. Two of the remaining PMS stars are found near the velocity of the Local Association together with a loose accumulation of some 5 or 6 further stars from the Pl_ZAMS age group. A relation of these stars with the Local Association may exist, but the errors and the scatter of the velocity vectors are quite large.
The W-V diagram displayed in the lower panel of Fig. 23 shows a similar trend in the distribution of the velocity vectors as in the U-V diagram, that is most PMS stars and many Pl_ZAMS stars are kinematically distinct from the Local Association.
Spectroscopic luminosity classification of the G-K stars based on the high resolution spectroscopy showed that 88% of the G-K stars are main-sequence stars or subgiants of luminosity classes V and IV, respectively. From IR photometric classification we concluded that all M stars are dwarf stars.
Significant lithium absorption lines were detected in a large fraction of stars with equivalent widths and abundances, respectively, above the level of the Hyades in about 50% of the stars. For the age distribution of the high-galactic latitude coronal sample this means that about half of the G-K stars are younger than the Hyades. About 25% of the G-K stars have an age comparable to that of the Pleiades, i.e. 100 Myr. A small fraction of less than 10% of the G-K stars is younger than the Pleiades. Most PMS stars, i.e. 8 out of 10, are located in area I. Only two PMS stars are found in area II and none in the remaining areas. This suggests a possible relation of the high-|b| PMS stars to the Gould Belt indicated in Fig. 1. However, the subsample formed by combining the stellar age groups PMS and Pl_ZAMS is spatially distributed in all directions covered by our study areas. At the same time half of its members show similar kinematical parameters independent of spatial location. This questions the relation to the Gould Belt. Rather, the space velocities suggest that these stars are members of a loose moving group with a mean velocity close to that of the Castor MG. For the Castor MG an age of Myr has been derived by Barrado y Navascués (1998). This would still be consistent with the Pl_ZAMS group. If some of the PMS stars are indeed kinematically related to the Castor MG this would indicate a large age spread in this moving group as they appear to be younger than 100 Myr, maybe even as young as 30 Myr.
We would like to thank the Deutsche Forschungsgemeinschaft for granting travel funds (Zi 420/3-1, 5-1, 6-1, 7-1). We further thank the staff at the German-Spanish Astronomical Centre, Calar Alto, in particular Santos Petraz, for carrying out part of the observing programme in service mode. 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. This research has made use of the SIMBAD and VIZIER databases, operated at CDS, Strasbourg, France.
Figure A.1 displays the measured lithium equivalent widths for each study area separately.
|Figure A.1: Equivalent widths of Li I as a function of plotted for each study area individually. The solid and dashed lines have the same meaning as in Fig. 16.|
In Table A.2 the basic parameters of the stellar sample are listed. The field and RASS names were taken from Paper III. Coordinates RA (2000) and Dec. (2000) of the optical counterparts are in succession either from Tycho-2, GSC-I, or GSC-II, whatever the source is for the V magnitude listed in column "V''. In column "Sp.type'' the revised spectral types with luminosity class are given for objects with new high resolution observations. Otherwise, spectral types from Paper III are given. Spectroscopic binaries are flagged by "SB2''. B160 is a triple system (SB3). The flux ratio is given for the revised V magnitudes and the RASS fluxes from Paper III. X-ray luminosities were calculated using the distances listed in column "dist''. The distances are flagged by "S'', "H'', "T'' or "I'', depending on whether they were derived from spectroscopy, Hipparcos, trigonometric parallaxes, or from the infrared colours, respectively. Distance estimates obtained by assuming luminosity class V are flagged by "M''. They should be considered as lower limits only.
Table A.3 lists the kinematical parameters. Heliocentric radial velocities and errors for single stars and for the primary component of spectroscopic binaries are given in columns and , respectively. For binaries columns and contain the heliocentric radial velocity and error of the secondary component, respectively. Proper motions and associated errors are listed in columns and for right ascension, and and for declination, respectively. The source catalog of the proper motions is denoted by respective flags: TY = Tycho-2, HI = Hipparcos, UC = UCAC2, US = USNO-B1.0, PP = PPM, ST = STARNET, TR = TRC, NL = NLPM1, CA = Carlsberg Meridian Catalogs. Also given are the galactic velocity components U, V, and W in the LSR frame with errors and , respectively. If the errors were larger than 30 km s-1 the space velocity components were omitted.
Table A.4 lists lithium data and rotational velocities. Equivalent widths of Li I are listed in column W(Li I ). Flags "h'', "l''or "m'' denote high-, low- or medium resolution measurements, respectively. Lithium abundances derived from W(Li I ) for the effective temperatures given in column are listed in column (Li). The last two columns list the rotational velocities, for single stars or primary components of binaries, and in the latter case for the secondary component.
|field||RASS name||RA||Dec||Sp. type||V||dist.|
|[km s-1]||[km s-1]||[mas yr-1]||[mas yr-1]||[km s-1]|
|field||Sp. type||W(Li I )||(Li)|
|[K]||[mÅ]||[km s-1]||[km s-1]|