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Astron. Astrophys. 325, 1115-1124 (1997)

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2. The data

2.1. Ca II H&K photometry

We have obtained Ca II H&K line photometry of 215 stars at the Mt. Wilson Observatory; most of the stars were observed within a few days of the scanning of these stars in the ROSAT All-Sky Survey. For only 9 stars the Ca II observations were separated by more than three weeks from the time of the All-Sky Survey observation.

The stars have been selected from the sample of Rutten (1987a), and are listed in Table 1. They are F-, G- and K-type stars of luminosity classes II to V, with known rotation rate. The sample stars are distributed over a large range in [FORMULA]  (0.4-1.5) and rotation period (1-400 days).

The Mt. Wilson H&K spectrophotometer measures the flux in two windows with 1Å or 2Å FWHM centered on the Ca II H&K line cores, and in two 20Å FWHM reference windows located on either side of the H&K doublet. The line-core emission index S is defined as the ratio between the number of counts in the line-core windows and the number of counts in the reference windows, scaled with a normalisation constant. A detailed description of the photometer and of the measurement procedure has been given by Vaughan et al. (1978). For most dwarfs and subgiants (luminosity classes IV to V) the 1Å FWHM passband ([FORMULA] -value) was used; for most giants and bright giants (luminosity classes II to III-IV) the 2Å FWHM passband ([FORMULA] -value) was used to accommodate their broader emission profiles in the H&K line cores (Wilson and Bappu 1957). Exceptions have been indicated in Table 1.

For most stars the Ca II H [FORMULA] K line-core emission index was measured two to six times, within an interval of a few minutes. The average S -values are listed in Table 1 (column 8). The listed uncertainty equals the standard deviation of the set of individual measurements; 84% of the measurements have uncertainties smaller than 2%. For a few stars only one measurement is available close to the X-ray observing time. For the relative uncertainty for these single measurement S -values we have taken 2%, somewhat above the mean relative uncertainty of 1.3% in our sample.

Fig. 1 shows a comparison with previous measurements of S -values, as listed by Rutten (1987a). The average spread is rather small, about 10%, although individual differences can occur of up to a factor 2. Relatively hot stars, with [FORMULA] (Fig. 1, top panel), show very little difference (reduced [FORMULA] of 0.69) between the measurements presented here and previously obtained measurements, suggesting that the amount of activity of these stars does not vary at a level exceeding the measurement uncertainty on time scales shorter than a few years. For the cooler stars (Fig. 1, bottom panel) the differences are on average much larger (reduced [FORMULA] of 9.5).

[FIGURE] Fig. 1. The [FORMULA] values derived here vs. the ones listed by Rutten (1987a). Top: [FORMULA] ; bottom: [FORMULA].

2.2. X-ray data

During the ROSAT All-Sky Survey the satellite scanned the sky in great circles perpendicular to the direction of the Sun. Any particular position on the sky was in the [FORMULA] field of view of the Position Sensitive Proportional Counter (PSPC) for about 30 seconds once every 90 minutes, during at least 2 days (depending on the ecliptic latitude). The PSPC is sensitive in the energy range 0.1-2.4 keV. For a detailed description of the satellite and the PSPC we refer to Trümper (1983) and Pfeffermann et al. (1988), and for a description of the All-Sky Survey to Voges (1992).

The X-ray count rates are derived as described in Chapter 2 of Piters (1995), and are given in Table 1 (column 9). We detected 134 X-ray sources out of the total of 215 stars, with the threshold value for detection set such that less than 0.5 false detections are expected. For the stars that were not detected, we derived a [FORMULA] upper limit from the total number of counts (as given by the Standard Analysis Software System, SASS; see Voges 1992, and Voges et al. 1992). These upper limits are also given in Table 1 (column 9).

There appears to be a systematic offset in the count rates determined in this paper and in a paper by Hempelmann et al. (1996); the latter are higher by about 30%. This difference is as yet not fully understood, but may be related to the exposure time corrections derived by the SASS, and used in the paper by Hempelmann et al. (1996). We stress that a constant normalisation factor that would have to be applied in case the offset is caused by an error on our part does not affect any of the conclusions reached in this paper, as it affects only the constant of proportionality in the fits.

The conversion of count rate to flux density at Earth [FORMULA] is given by

[EQUATION]

where [FORMULA] is the energy-conversion factor, derived from the ROSAT hardness ratio h and from the hydrogen column density [FORMULA], following the method described in Piters (1995; Ch. 2). The hardness ratio and its uncertainty are listed in column 10 of Table 1. For nearby stars in the galactic plane (distance less than 200 pc and galactic latitude between [FORMULA] and [FORMULA]) we derived [FORMULA] from Paresce (1984), while for more distant stars we estimated [FORMULA] from the interstellar reddening [FORMULA] using the expression [FORMULA] (Bohlin et al. 1978). The spread around this relationship is about 30%. The adopted [FORMULA] values are listed in Table 1, column 11. The distance is derived from the parallax or, if the parallax is not known, from the distance modulus, using the absolute magnitudes listed by Schmidt-Kaler (1982).

The ROSAT hardness ratio used here is defined as the ratio between the source count rate in PSPC channels 41-240 ([FORMULA] 0.4-2.4 keV) and the total source count rate. The hardness of the soft X-ray spectrum is a measure for the mean coronal temperature. The hardness ratio increases with temperature up to 5 MK, and then decreases slightly for higher temperatures (see Chapter 2 of Piters, 1995). For spectra with only a few counts, this hardness ratio can still yield valuable information about the coronal temperature structure, provided that the column density is known: for high values of the column density the number of counts in the low-energy band is suppressed, and consequently the hardness ratio is higher.

For the main-sequence stars in our sample we see a strong correlation of hardness ratio with the X-ray surface flux density (Fig. 2, top; the derivation of the surface flux density is described in the next section), suggesting (see Schrijver et al., 1987) that as a star becomes more active, it will either heat up the coronal material as a whole or produce more high-temperature plasma. Both options have the effect of increasing the hardness ratio of the spectrum. For giants this trend is somewhat less pronounced (Fig. 2, bottom).

[FIGURE] Fig. 2. Hardness ratio as a function of X-ray surface flux density for main-sequence stars (LC IV-V and V) (top) and giants (LC IV and up) (bottom). Stars with hydrogen column density [FORMULA] are indicated by small circles. The average uncertainty is indicated by the cross in the upper right corner.

Note that since the countrate-to-flux conversion factor [FORMULA] depends on the hardness ratio, it depends on the X-ray flux density itself! Not taking into account this dependence (for simplicity, the coronal temperature structure is usually assumed to be the same for all stars) would therefore affect the slope of the flux-flux relationships, as discussed in Section 5.2.

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© European Southern Observatory (ESO) 1997

Online publication: April 28, 1998

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