q 2000 RAS

Simultaneous ROSAT XRT and WFC observations of a sample of activedwarf stars

V. Tsikoudi,1 B. J. Kellett2w and J. H. M. M. Schmitt31Laboratory of Astronomy, Department of Physics, University of Ioannina, Ioannina, Greece2Space Science Department, Rutherford Appleton Laboratory, Chilton OX11 0QX3MPI fuÈr Extraterretriche Physik, Garching, Germany

Accepted 2000 July 27. Received 2000 July 24; in original form 1999 July 12

A B S T R A C T

The X-ray observations of the ROSAT-PSPC All-Sky Survey have revealed bright and

energetic coronae for a number of late-type main-sequence stars, many of them flare stars.

We have detected 31 X-ray flares on 14 stars. A search for simultaneous X-ray and EUV

(extreme ultraviolet) flares using ROSAT Wide Field Camera survey data revealed a large

number of simultaneous flares. These results indicate that the heating mechanisms of the

X-ray and EUV-emitting regions of the stellar coronae are similar. We find X-ray quiescent

variability for nine of the 14 stars and simultaneous X-ray and EUV quiescent variability for

seven of these nine stars. These results imply that the stellar coronae are in a continuous state

of low-level activity. There are tight linear correlations of X-ray flare luminosity with the

`quiescent' X-ray as well as with the stellar bolometric luminosity. The similarity between

the X-ray-to-EUV quiescent and flare luminosity ratios suggests that the two underlying

spectra are also similar. Both are indeed consistent with the previously determined Einstein

two-temperature models. We suggest that both the variability and spectral results could

indicate that the quiescent emission is composed of a multitude of unresolved flares.

Key words: stars: activity ± stars: flare ± stars: late-type.

1 I N T R O D U C T I O N

The coronal X-ray properties of cool main-sequence stars have

been investigated in the past by Einstein and EXOSAT observa-

tions (Vaiana et al. 1981; Ambruster, Sciortino & Golub 1987;

Pallavicini, Tagliaferri & Stella 1990, and others). For a review on

this subject see Haisch (1983) and references therein.

More recently the ROSAT All-Sky Survey (RASS), carried out

in the second half of 1990 and the very early part of 1991,

collected an almost complete set of observations in both the X-ray

(6±120 AÊ ) band and, simultaneously, in the Extreme Ultraviolet

(EUV; 60±200 AÊ ) in two bands. Several investigators have

analysed the ROSAT X-ray Telescope (XRT) and ROSAT Wide

Field Camera (WFC) EUV data for different classes of stars

(Haisch & Schmitt 1994; Schmitt 1994; Schmitt, Fleming &

Giampapa 1995; and others). The high sensitivity of the

instruments aboard ROSAT allows the study of the quiescent

emission of a large number of stars. In addition to the X-ray

emission, the EUV emission of stars extends the total band pass of

the observations and opens up cooler regions of the stellar

atmosphere. This allows for an indepth study of the physical

conditions and processes of stellar coronae.

In previous work we have performed a systematic study of

coronal activity and variability for a large number of main-

sequence stars using the ROSAT-WFC observations (Kellett &

Tsikoudi 1997; Tsikoudi & Kellett 1997). Our sample included

127 optically selected, photospherically and chromospherically

active stars ranging from the low-activity dF5±dK5 to the very

active dK5e±dM6e flare stars. We have detected quiescent EUV-

radiation on 47 late-type stars. A large number of these were also

observed to flare and/or to exhibit low-level variability, i.e.

excluding any obvious flare events. These variations were on time-

scales of several hours to several days. The flare events covered

over two orders of magnitude in flux, and the low-level variations

probably indicate the presence of many undetected flares, which

might be the cause of the continuous heating of the stellar

coronae.

One approach to investigate the possible process (or processes)

of coronal heating is the intercomparison of coronal emission

from different temperature levels of the coronae of stars with

similar photospheres. As we will demonstrate below, the WFC S1

and S2 filters are much more sensitive to temperatures below

106 K than the XRT1PSPC. Our present sample of late-type main-

sequence stars offers such an opportunity. The RASS observations

enable us to investigate the simultaneous X-ray and EUV stellar

emission and to perform a systematic search for X-ray and EUV

flares of a large number of stars. The goal of this present work was

to search the ROSAT XRT data base at the positions of 14 of the

Mon. Not. R. Astron. Soc. 319, 1136±1146 (2000)

w E-mail: b.j.kellett@rl.ac.uk

main-sequence stars for which we already had positive EUV

detections (Tsikoudi & Kellett 1997). Most of these stars are well-

known flare stars.

In the sections that follow we report on the observations, both in

X-rays and in the EUV (Section 2), and then in Sections 3 and 4

we describe the results obtained. The discussion and conclusions

of this study are presented in Section 5.

2 X - R AY A N D E U V O B S E RVAT I O N S

The ROSAT satellite was launched in 1990 June. It carried the

X-ray telescope which, for the All-Sky survey, utilized a Position

Sensitive Proportional Counter (PSPC) detector (Pfeffermann,

Briel & Hippmann 1986; Trumper et al. 1991). The ROSAT

satellite also carried a second instrument, the UK Wide Field

Camera, a co-aligned EUV telescope (Sims et al. 1990). The

satellite surveyed over 92 per cent of the sky in both the X-ray and

EUV bands (the EUV survey was performed in two overlapping

bands). The sky survey occupied the first six months of the

mission and was carried out following great circles that passed

through the ecliptic poles, in a plane approximately perpendicular

to the solar vector. Each great circle survey scan lasted one

satellite orbit (approximately 96 min), and advanced at the rate of

one degree per day, to cancel the Earth's motion around the Sun.

The WFC field of view is circular with a diameter of five

degrees; a typical source was therefore observed every 96 min, for

a period of between about 30 and 80 s, over a total interval of

about 5 d. The WFC used two filters, referred to as S1 and S2 (`S'

for survey), which were alternated once per day. The band passes

defined by the two filters were 60±140 AÊ (S1) and 110±200 AÊ

(S2). The S2 filter has slightly the larger effective area of the two

filters. However, it also suffers from a higher background count

rate. For more details of the WFC all-sky survey analysis methods

see Pounds et al. (1993) and Pye et al. (1995).

The PSPC detector has a two-degree field of view; so, a

particular source would have been scanned for a minimum of 2 d.

The energy range of the PSPC is roughly 0.1±2.4 keV or in the

waveband region of 5±124 AÊ , and the detector also possesses a

modest spectral capability which will not be considered here. As

can be noted the PSPC response overlaps, to a certain extent, with

that of the WFC1S1 band. In this overlap region, the PSPC

reaches ,10 times the effective area of the WFC1S1. However,

the WFC survey coverage for any given source was over twice that

of the X-ray telescope.

We selected a subset of 14 main-sequence stars from our full set

of EUV detections presented previously (Tsikoudi & Kellett

1997). The stars we selected for this study were required to cover

a wide range of spectral types and activity levels, and which we

knew had been detected in the EUV in both quiescence and flaring

states. As expected, there was quiescent X-ray emission observed

from all the stars in our sample. There were also at least one flare

observed on each of the stars in our sample; 11 of these 14 stars

are previously known (optical) flare stars.

We applied the appropriate corrections for vignetting as well as

instrumental corrections to the observations but we did not correct

for background emission since it is much less than the X-ray

signal, the background contribution is always much less than the

statistical error for each data point (Haisch & Schmitt 1994).

We proceeded to obtain light curves for the XRT data to

compare with our EUV data. From these data, we can then

estimate count rates both outside any obvious flares and also for

the flares themselves. Flare orbits were defined by an automatic

procedure explained in detail in Section 4. Basically, flares are

defined as extreme positive outliers from the `mean' level

(typically 3±10s above). The actual definition used was .1.8s(see Table 3). All the photon-counts outside the flare events were

used to estimate the mean count rates of the quiescent state. The

mean count rates were calculated, for each star, over the entire

time interval of the ROSAT±PSPC coverage. We also obtained the

counts per second at the peak of each flare observed. The count

rates are given in Table 1, together with other relevant information

concerning the observations both in the X-ray and in the EUV

bands. The name of the star (its Gliese or SAO catalogue number

and common name) is given in columns 1 and 2. In columns 3 and

4 we give the time coverage of the XRT survey observation (the

start and end times in Julian days). In columns 5±7 we give the

XRT/PSPC information. The mean count rates of the quiescent

state are given in column 5, the time of occurrence of each flare, if

any, is given in column 6 and the count rate at the maximum

(observed) of the flare is given in column 7. In columns 8±11 we

give the EUV observation details. The mean count rates outside

the flares are given in column 8, the time of occurrence of the

flare(s) is given in column 9, the count rate during the flare

maximum is given in column 10 and the filter in which the flare

was observed is given in column 11. Column 12 notes any

differences between the time of the X-ray flare and that of the

EUV flare (i.e. notes whether the EUV event is before or after the

X-ray event ± the simultaneous events are kept blank in this

column).

2.1 Relative X-ray and EUV sensitivity and temperature

response

Before we go on to describe the results in detail it is worthwhile

discussing the relative performances and sensitivities of the two

EUV survey filters (S1 and S2) and the XRT/PSPC. In Fig. 1 we

show the ratio of S1/S2 and also each WFC filter relative to the

PSPC for both single-temperature and two-temperature plasma

emission models. We have used the Landi & Landini (1998)

plasma models for these simulations. The various response curves

were taken from the data archive at MPI in Garching (PSPCC) or

from WFC calibration files in the UK. Fig. 1(a) shows the single-

temperature results. The WFC survey filter ratio shows a broad

peak for temperatures of several millions K where the two filters

have their combined maximum sensitivities. The EUV-to-PSPC

ratios show quite large variations with temperature. In Fig. 1(b) we

show the various filter ratios for the two-temperature simulations.

Here we show two cases: an equal emission measure (EM) case

(solid lines) and also the case of the higher temperature

component being twice the emission measure of the cooler

component. The assumed temperature difference between the two

components is 0.8 in log temperature. This temperature difference

and the larger EM for the hotter component are both consistent

with Einstein Observatory results for dMe stars (Schmitt et al.

1990). These two-temperature simulations show a broad peak for

the WFC filter ratio for temperatures in the few tens of millions K

range (which is exactly where Einstein results put the hotter

component) and a relatively flat S1-to-PSPC ratio for hotter

component temperature in above about 1 million K. The S2/PSPC

ratio does show a drop in sensitivity for temperature of 2±5 �107 K (which leads to the peak in the WFC filter ratio already

mentioned above).

XRT and WFC observations of activity 1137

q 2000 RAS, MNRAS 319, 1136±1146

Table 1. XRT/WFC flare stars ± data.

Star Gl Name XRT survey XRT/PSPC observations EUV/WFC observations Commentsb

Start JDa End JDa Mean XRT counts s21 Flare JD Flare counts s21 Mean EUV counts ks21 Flare JD Flare counts ks21 Filter

(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12)54.1 YZ Cet 085.80 087.94 0.220^ 0.026 086.81 0.636^ 0.203 14.09^ 13.94 ,50.0 S1

087.87 6.322^ 0.954 087.87 132.1^ 135.9 S265 UV Cet 254.22 256.30 1.071^ 0.059 254.30 1.957^ 0.661 21.34^ 4.37 254.30 73.3^ 36.0 S1

255.96 1.259^ 0.393 255.96 22.0^ 26.3 S1171.2 283750 125.62 127.68 2.639^ 0.081 126.88 1.982^ 0.462 48.07^ 7.13 126.82 79.9^ 46.2 S2 EUV , X233 OU Gem 151.92 153.92 1.710^ 0.244 152.26 2.204^ 0.569 43.62^ 19.09 152.39 18.6^ 37.4 S1 EUV . X278c YY Gem 166.66 168.74 2.840^ 0.115 167.93 8.890^ 0.720 55.80^ 8.35 167.93 179.1^ 56.1 S1

168.40 5.151^ 0.760 168.40 54.0^ 48.5 S2388 AD Leo 207.40 209.33 3.000^ 0.106 208.60 2.508^ 0.506 50.27^ 8.43 208.53 53.0^ 45.3 S2 EUV , X

209.10 2.603^ 0.693 209.00 37.8^ 41.6 S2 EUV , X398 RY Sex 216.80 218.70 0.324^ 0.032 218.20 0.481^ 0.210 10.35^ 4.60 218.34 12.1^ 23.8 S1 EUV . X551 Prox. Cen 111.93 114.80 1.066^ 0.064 112.80 2.800^ 0.442 22.29^ 8.04 112.80 111.5^ 45.8 S2

113.46 3.091^ 0.419 113.33 45.5^ 32.5 S1 EUV , X114.53 1.549^ 0.507 114.53 55.6^ 52.3 S2

644 Wolf 630 127.72 129.85 3.253^ 0.118 128.58 2.392^ 0.505 73.08^ 9.30 128.58 133.4^ 55.8 S2128.98 2.525^ 0.517 129.12 10.5^ 35.9 S1 EUV . X129.52c 9.956^ 0.989 129.72 107.7^ 51.3 S1 EUV . X

719 BY Dra 163.58 171.32 2.649^ 0.054 166.12c 7.023^ 0.702 52.05^ 3.57 166.12 93.9^ 42.6 S2167.32 1.951^ 0.432 167.25 40.1^ 36.2 S1 EUV , X167.72 2.122^ 0.448 167.65 22.9^ 36.3 S1 EUV , X169.05c 1.998^ 0.485 169.05 76.6^ 42.5 S1

799 AT Mic 181.58 183.28 3.693^ 0.137 182.31 2.151^ 0.506 51.64^ 15.03 182.45 108.0^ 61.5 S2 EUV . X182.65 1.410^ 0.493 182.65 58.6^ 50.0 S2

803 AU Mic 182.38 184.41 3.744^ 0.143 183.21 32.067^ 1.296 79.39^ 10.38 183.21 302.7^ 73.5 S1183.61 4.713^ 0.638 183.75 106.1^ 53.9 S1 EUV . X183.88 6.657^ 0.766 183.81 102.4^ 54.6 S1 EUV , X

873 EV Lac 244.07 246.88 2.823^ 0.147 244.81 3.694^ 0.636 68.61^ 8.02 244.81 62.6^ 45.1 S2245.68c 8.027^ 0.691 245.74 60.3^ 43.9 S1 EUV . X246.21 22.081^ 1.161 246.21 230.7^ 61.3 S1

86592d V775 Her 164.58 167.44 2.000^ 0.064 166.31 5.401^ 0.576 46.86^ 6.14 166.31 158.8^ 55.2 S2167.24 2.172^ 0.698 167.31 5.7^ 29.5 S1 EUV . X

a JD ± given as JD 244 8000.0b Comments: EUV , X ± the EUV flare peak is BEFORE the X-ray; EUV . X; EUV peak AFTER X-ray peak.c The flare is extended over several ROSAT satellite orbits or is a multi-flare event. (Only the X-ray flare peak is given.)d This is an SAO identification number.

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So, from the above observations we can conclude that the S1

and PSPC are likely to respond in relatively the same way for

`realistic' two-temperature quiescent stellar spectra, but that the

two instruments could respond differently to single-temperature

`flare' emission spectra. Since we know that hot flare plasma is

produced by some form of magnetic heating mechanism in flare

loops, it is possible that similarities (or differences) between the

responses of the XRT1PSPC and the WFC/EUV to flare emission

and the more constant quiescent emission could help us in

understanding more about the mechanism that generates and

maintains the hot coronae in active dwarf stars.

3 L I G H T C U RV E S A N D F L A R E E V E N T S

In Figs 2 and 3 we show examples of X-ray and EUV light curves;

they are for the stars Gl 551 and SAO 86592. The light curves

cover the time interval of the XRT/PSPC observation-coverage

and is given in both JD and ut. As we notice, there are several

flare events which occur in both the X-ray and EUV wavebands.

Most of the X-ray flares have their EUV counterpart. However, for

some X-ray flares, there is no obvious EUV counterpart. Similar

results can also be found for many of the other stars in our sample.

However, in all cases there is a possible EUV excess within one or

two satellite orbital scans (i.e. ^3.2 h) of the XRT event. In

Section 5 we will estimate the statistical likelihood that these EUV

events are random unrelated flares. These results are summarized

in Table 1 where we record all the flares found in the XRT data

(columns 6 and 7), and we give the EUV count rates (columns 9

and 10) of the closest EUV excess to the XRT event. In the EUV

data, the `flare' event corresponding to a clear XRT excess was

very often less than 3s significance and in some cases no clear

excess was visible at all (i.e. X-ray only flares). We should also

note that the values recorded are only for the `peak' orbit/scan. If

the flare event extended over more than one orbit, the `rise/decay'

orbits are not included in the count rates given.

For stars like Gl 54.1, 65, 171.2, 233, 388 and 398, the EUV

flares are only marginal, whereas the X-ray flares are well defined

and above the 3s level. For Gl 278c, 551, 644, 719, 799, 803, 873

and SAO 86592, the EUV flares corresponding to the X-ray flares

are significant. For the star SAO 86592, we observe a large flare at

JD , 244 8166:3; in both the X-rays and EUV; a smaller X-ray

flare at JD , 244 8167:25 has no apparent EUV counterpart

(Fig. 3). Taking the peak parameters for the first flare we have an

X-ray/EUV ratio of 34:1. This is consistent with an S2/PSPC

single-temperature plasma above about 4 � 107 K or between 3

and 7 million K [Fig. 1(a)]. For the same flux ratio, the second

flare should have produced about 60 c ks21 in the WFC ± some

10� more than actually seen. Either the flare was extremely hot

(above 108 K) or the EUV emission was somehow delayed (or

absorbed?). Similar examples can be found for many of the other

stars in our sample, as can be judged from Table 1. We have also

noted in Table 1 whether the EUV flare peaks occur before or after

the X-ray peak (although we only considered orbits within ^2 of

the X-ray orbit). Given the relatively low frequency of flares in our

sample, it is perhaps unlikely that these non-simultaneous flares

are separate flares. However, we note that Table 1 lists that very

nearly 50 per cent of all the flares do not have a simultaneous

EUV response. We will return to this point in Section 5.

4 X - R AY Q U I E S C E N T E M I S S I O N A N D

Q U I E S C E N T VA R I A B I L I T Y

The observed count rates of the quiescent state (Table 1, column 5)

and of the flares (Table 1, column 7), were converted to X-ray

Figure 1. (a) The relative temperature sensitivity of the WFC survey filters with respect to each other and to the XRT1PSPC for single temperature coronal

models. We have used the plasma models of Landi & Landini (1998), assuming an electron density of 1012 cm23. (b) As (a) but for two-temperature coronal

models. Here we have calculated the ratios for two cases. Solid lines represent the example of equal emission measures for the two components and dashed

lines show the case of the hot emission measure being double the cool one. The cool component was taken to be 0.8 less than the hotter component in log

temperature. The dotted line is the single-temperature WFC filter ratio from (a).

XRT and WFC observations of activity 1139

q 2000 RAS, MNRAS 319, 1136±1146

luminosities using the conversion factor 6 � 10212 erg cm22 s21

per XRT/PSPC count s21 (Haisch & Schmitt 1994) and distances

from Gliese & Jahreiss (1991). We have also converted the EUV

count rates of the quiescent state and of the flares (Table 1,

columns 8 and 10) to EUV luminosities using the single con-

version factor of 4 � 10211 erg cm22 ks21 per WFC count ks21 for

both filters. This is equivalent to assuming a particular

temperature range in which both filters are equally sensitive. We

will see below that this assumption is consistent with the WFC

observations for all the stars in the present sample.

In Table 2 we give in the first two columns the star name and

distance in parsecs used in calculating the luminosities for the

quiescent X-ray and EUV (columns 3 and 4) and then for each

flare (X-ray, EUV; columns 6 and 7). We also show the bolometric

luminosities of the stars (column 5), which were taken from

Kellett & Tsikoudi (1997), and are calculated from the absolute

magnitudes of each star. The flare as well as the quiescent X-ray

luminosities of these active stars vary by a factor of almost 1000

from one star to another although these stars cover a very narrow

spectral range (dK5±dM5). We note that the EUV non-flare

luminosity varies by a similar factor to the XRT.

We compare the X-ray quiescent luminosities to the EUV and

bolometric luminosities in Figs 4(a) and (c), respectively. There is

a good relationship between the stellar X-ray luminosity, LX, with

both the EUV and bolometric luminosities, LEUV and LBol. As can

be seen in Fig. 4(c) (and Table 2, column 6), the LX/LBol ratio for

Figure 2. The XRT and WFC light-curves for Proxima Cen. (Gl 551). This shows a `typical' example of the survey data. In the WFC panel, the diamond

symbols represent S1 data, while the squares are S2 data. The data gaps in this light curve are caused by passages through the Earth's radiation belts. The dot-

dash lines indicate the calculated quiescent mean level.

Figure 3. The XRT and WFC light curves for SAO 86592. For the definition of symbols, see Fig. 2.

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Table 2. XRT/WFC flare stars ± luminosities.

Star (Gl or SAO) d (pc) LXRT (erg s21) LEUV (erg s21) LBol (erg s21) log LXRT

LBolLXRT (flare) (erg s21) LEUV (flare) (erg s21) WFC EUV full survey

(quiescent/mean) NS1 NS2 Filter ratio

(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)54.1 3.7 2.21� 1027 9.55� 1026 1.26� 1031 23.756 6.39� 1027 ,3.39� 1027 (No quiescent S2A flux)

6.35� 1028 8.95� 1027

65 2.6 5.31� 1027 9.93� 1026 6.17� 1030 23.065 9.69� 1027 3.41� 1027 23 21 1.26^ 0.466.24� 1027 1.02� 1027

171.2 16.4 5.07� 1029 6.72� 1028 5.62� 1032 23.045 3.81� 1029 1.12� 1029 26 22 1.11^ 0.25233 15.0 2.77� 1029 5.39� 1028 1.29� 1033 23.668 3.57� 1029 2.30� 1028 25 10 0.85^ 0.44278c 14.6 4.37� 1029 6.74� 1028 3.47� 1032 22.899 1.37� 1030 2.16� 1029 18 21 1.00^ 0.25

7.93� 1029 6.53� 1028

388 4.9 5.18� 1028 7.35� 1027 7.59� 1031 23.166 4.33� 1028 7.75� 1027 14 25 0.88^ 0.224.49� 1028 5.53� 1027

398 13.7 4.36� 1028 1.20� 1028 4.47� 1031 23.010 6.48� 1028 1.41� 1028 11 14 1.63^ 0.89551 1.3 1.28� 1027 1.90� 1026 6.17� 1030 23.681 3.37� 1027 9.52� 1026 17 25 0.86^ 0.41

3.73� 1027 3.88� 1026

1.87� 1027 4.75� 1026

644 6.5 9.86� 1028 1.66� 1028 9.12� 1031 22.966 7.25� 1028 2.95� 1028 20 21 1.00^ 0.207.65� 1028 2.32� 1027

3.02� 1029 2.38� 1028

719 17.2 5.65� 1029 8.70� 1028 7.24� 1032 23.108 1.50� 1030 1.57� 1029 77 74 1.08^ 0.104.16� 1029 6.70� 1028

4.53� 1029 3.83� 1028

4.26� 1029 1.28� 1029

799 8.1 1.76� 1029 1.98� 1028 5.89� 1031 22.525 1.02� 1029 4.14� 1028 11 17 1.05^ 0.356.71� 1028 2.25� 1028

803 9.4 2.35� 1029 4.03� 1028 2.34� 1032 22.999 2.01� 1030 1.54� 1029 13 13 0.92^ 0.212.96� 1029 5.39� 1028

4.18� 1029 5.20� 1028

873 5.1 5.22� 1028 1.15� 1028 5.01� 1031 22.982 6.83� 1028 1.05� 1028 21 28 1.01^ 0.161.48� 1029 1.00� 1028

4.08� 1029 3.87� 1028

86592 24.4 8.54� 1029 1.44� 1029 3.39� 1034 24.598 2.31� 1030 4.89� 1029 24 36 1.19^ 0.249.28� 1029 1.76� 1028

Notes: NS1 & NS2 are the number of good S1 & S2 orbits used to calculate the mean filter countrates and hence the filter ratio. The time interval in the WFC survey over whichthese orbits were accumulated was defined by the shorter XRT survey interval.

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the majority of our stars cluster around the 1023 so-called

`saturation' limit found in many types of active stars from dMe

through to evolved stars in RS CVn binaries (e.g. Haisch &

Schmitt 1994). The X-ray emission increases as the EUV emission

increases, as expected, and also as the bolometric luminosity

increases. A similar relation was found between LEUV and LBol by

Kellett & Tsikoudi (1997). Table 2 (column 11) lists the WFC

survey filter ratio (S2/S1), and this is also shown in Fig. 4(b). This

shows that we can `merge' the two WFC filters to produce a

`combined' EUV measurement, since for all the sources here the

filter ratio is consistent with unity. From Fig. 1(b), this filter ratio

suggests a two-temperature spectrum with the hotter component

around 2±5 � 107 K: This would be perfectly consistent with the

earlier findings from X-ray spectra (Schmitt et al. 1990).

Comparison of the X-ray flare luminosity with the X-ray

quiescent (non-flare) luminosity is shown in Fig. 5. Again, we see

a tight correlation between the two, with the more luminous flares

detected on the more luminous (mean/quiescent) stars. This is

likely to be a selection effect, given that we have to be able to

`detect' the flare as a distinct event above any quiescent emission.

Fig. 5 does illustrates that above the observed `base' flare flux, we

have also detected several larger flares ± these are on the stars Gl

54.1, Gl 803 and Gl 873. In particular, the flare on Gl 54.1 is

impressive, since this is one of the least X-ray luminous stars in

our sample [only Proxima Cen. (Gl 551) is less luminous]. In fact,

the flare peak is very nearly 30 times the mean/quiescent count

rate. (The flare is essentially a single orbit event, the star is very

nearly back to quiescent in the next orbit. Unfortunately, these

were the last two orbits of XRT survey data on this star.)

As can be seen from Table 1, of the 31 flares listed, one is

completely undetectable in the EUV, 15 flares produce a

simultaneous response and for the remaining 15 flares the EUV

response is apparently in a different orbital scan or completely

unrelated. Fig. 6, which displays the X-ray and EUV responses for

each flare, does show a relatively good correlation even though we

have included the 15 non-simultaneous events. Just taking the 15

simultaneous flares, the EUV-to-X-ray ratios are in the range of

0.015±0.05, consistent with temperature in either the 2±8 million

K or , 2 � 107 K range. The Pearson correlation coefficient for all

30 flares is 0.927 and the Spearman Rank coefficient is 0.887.

Both of these statistics confirm the good correlation. However, we

would caution that although this plot covers over three orders of

magnitude in luminosity, it is based on a `dynamic' range of only

60 times in XRT count rates and a similar range in S1/EUV. We

note that the trend line we show in Fig. 4(a), relating the quiescent

X-ray and EUV luminosities, is a surprisingly good match to the

Figure 4. The XRT `mean' luminosity versus: (a) The `corrected/

simultaneous' EUV luminosity. These estimated quiescent fluxes exclude

any obvious flare events, and the EUV data are restricted to the time range

appropriate for the XRT data. (b) The EUV survey filter ratio (S2/S1). This

is to show that the two EUV survey filters are essentially equivalent for the

spectra emitted by these active stars. (c) The total bolometric luminosity.

The solid and dashed lines indicate constant LX/LBol ratios of 1023 and

1024, respectively.

Figure 5. For all the 31 X-ray flares listed in Table 1, we plot the peak flare luminosity against the `quiescent' X-ray luminosity. The quiescent flux of each

star is based on the mean underlying flux of each star, ignoring any obvious flares. The peak flare luminosity is then taken as the maximum observed flux

above this mean level; i.e. a flare-only flux.

1142 V. Tsikoudi et al.

q 2000 RAS, MNRAS 319, 1136±1146

flare data. If the non-simultaneous EUV flares were just `back-

ground' events, and unconnected with the observed X-ray events,

it is perhaps unlikely that they would `know' anything about the

X-ray flux, and, therefore, these flares should be expected to

`scatter' more than is apparently the case. We will return to this

point in Section 5.

4.1 X-ray variability

Next we tested for X-ray variability during the `quiet state' (i.e.

outside the obvious flares) of each star. We estimated the X-ray

variability and calculated X-ray `quiet-state means' (these are the

values already discussed in Table 1 and various figures). The

method we used to estimate these means is discussed in detail in

Haisch & Schmitt (1994). It is a non-trivial problem to distinguish

the intrinsic variability of X-ray sources observed in the RASS,

since the PSPC detector itself induces a certain variability due to

the window support ribs. Haisch and Schmitt describe a method of

taking account of this `detector intrinsic variability' and thereby

allowing the true source variability to be assessed. So, we have

followed their `2 step' method. However, we have added an

additional step ± a 0th step ± to first take account of the very

obvious `flare scans' present for all the stars in our sample.

The full method we adopted was the following. We rejected the

very short scans with durations of less than 10 s. These are the

BAD scans which only sample the extreme limit of the circular

PSPC detector. All the remaining scans result in the GOOD figure

quoted in Table 3. Our additional 0th step comes next. We obtain

the overall mean of all the GOOD scans and then eliminate any

scans that are more than 1.8s above this mean. The factor of 1.8sis somewhat arbitrary, but the results are not very sensitive to the

exact value since the `flare' scans that we are attempting to

exclude are typically many (3±10)s above the mean. If any orbits

were excluded, we recalculated the mean and tested again for

orbits 1.8s above the mean, continuing until no more points are

found. The total number of scans excluded by this cut are recorded

as the FLARE count in Table 3. We then continued with the

method of Haisch & Schmitt (1994). Step 1 is now to find the

mean of all the GOOD±FLARE scans, and reject the 1.8s low

scans. These are attributed to the detector artifacts and labeled

LOW in Table 3. Step 2 is then used to determine the number of

scans that lie 1.8s above the mean. These are the HIGH scans and

they are attributed to true variations in the star itself. The choice of

1.8s used in steps 1 and 2 is discussed in detail by Haisch &

Schmitt (1994). Having excluded the FLARE, LOW and HIGH

scans, we can finally calculate the mean of the remaining scans. It

is this value that we have taken to represent the mean level of our

sample ± a `quiet-state mean' value which attempts to exclude

flares and other positive excesses and at the same time also

excludes possible detector artifacts.

Table 3 summarizes all the results of the variability analysis for

our sample of 14 stars. As discussed by Haisch & Schmitt (1994),

for a 1.8s error and a Gaussian error distribution, 92.8 per cent of

the data points should lie within 1.8s of the mean. Alternatively,

this can be expressed in terms of the number expected to exceed

(on average) the mean by 1.8s ± which is one scan in every 28.

The global statistic for our sample is as follows. We have a total of

501 good scans and 60 of these were judged to be LOW. This is

Figure 6. For the 30 X-ray flares where we also see an EUV response listed in Table 1, we plot the peak X-ray flare luminosity against the peak EUV flare

luminosity. These peak estimates subtract the quiescent flux of each star. The trend line from Fig. 3(a) is also shown (exactly), and represents a remarkably

good representation of the flare data. (The 15 flares that peak simultaneously in both the X-ray and EUV bands are indicated with a filled box symbol.)

Table 3. XRT flare stars ± variability.

Star GOOD FLARE HIGH LOW Var?a WFC

54.1 32 2 1 2 V/-65 32 3 0 6 C/C

171.2 31 1 0 2 C/C233 28 2 0 0 C/V278c 30 5 4 6 4:19 V/V388 30 4 4 3 4:23 C/V398 23 1 0 1 C/C551 29 6 2 4 2:19 C/C644 30 6 4 3 4:21 V?/V719 108 14 17 15 17:79 C/C799 22 4 4 3 4:15 V?/V803 28 4 6 5 6:19 V/C873 35 10 5 7 5:18 V?/V

86592 43 3 2 3 2:37 V?/C

a Variability is judged on the basis of the ratio HIGH:(GOOD±FLARE±LOW). If this exceeds the value expected (Gaussian) rateof 1:28, then we suspect `quiescent' variability. See text for fulldetails of the method.

XRT and WFC observations of activity 1143

q 2000 RAS, MNRAS 319, 1136±1146

about 12.0 per cent of the total number of scans (Haisch and

Schmitt found 15.4 per cent for their sample). Given the different

nature of the two samples, these figures seem to be in reasonable

agreement. We also find that a total of 65 scans are FLARE and

another 49 scans are HIGH (13.0 per cent and 9.8 per cent,

respectively). This is substantially higher than what Haisch and

Schmitt found in their sample, but this can easily be explained by

the fact that our sample is predominantly composed of well-

known flare stars, while their sample was composed of active late-

type giants and RS CVn systems which can be expected to flare

rather less often than the dKe and dMe stars in our sample. Given

that we actually recorded 31 `flare events', it follows that the

`typical flare' covered approximately two to three orbits/scans (if

we add the HIGH scans to the FLARE scans). This gives a mean

duration of 1.5±3 h for the X-ray flares. The last column in Table

3 records whether we found the star in question to be variable

(`V') or constant (`C') in the S1 and S2 WFC filters in our earlier

analysis (Tsikoudi & Kellett 1997). `V?' indicates sources that we

thought were `possibly variable'

The above analysis revealed significant X-ray quiescent

variability for nine of the 14 sample stars. These stars are: Gl

278, 388, 551, 644, 719, 799, 803, 873 and SAO 86592. This is in

general agreement with the earlier analysis of the EUV data

(Tsikoudi & Kellett 1997), which found strong evidence for

variability in both filters from Gl 278, 644, 799, and 873 and

evidence for variability in a single filter from Gl 54.1, 233, 388,

803 and SAO 86592. Therefore, the only differences between the

two wavebands are that the WFC analysis suspected variability on

Gl 54.1 (which is an extremely weak source in the WFC data,

except for the obvious flare events) and also on Gl 233; and,

secondly, the stars Gl 551 and 719, which were judged `constant'

(i.e. excluding the obvious flares). Gl 551 shows the second

lowest `variability' index in the XRT data, at 10.5 per cent (the

lowest is SAO 86592 at 5.4 per cent, where the threshold is

3.6 per cent).

The fact that we have observed variability in X-ray and

EUV bands for seven out of the nine stars variable in the XRT

alone certainly enhances the probability that this simultaneous

quiescent variability is real and therefore significantly strength-

ens the earlier results, which were based solely on the EUV

observations.

5 D I S C U S S I O N A N D C O N C L U S I O N S

The sensitivity of the ROSAT±PSPC detector allows weaker flare

events to be more easily identified than previous X-ray instru-

ments (Einstein, EXOSAT, etc.). In addition, our analysis of the

ROSAT±WFC survey observations (Tsikoudi & Kellett 1997) has

proven successful in detecting EUV emission from faint stars. Our

study of a large number of main-sequence stars, of spectral types

between F5 and M6, indicated that the stars flare frequently at the

EUV and that there is significant small-scale flaring on time scales

of hours and days, which might very well be the cause of the

`continuous' heating necessary to generate the stellar coronae in

these active stars (Kellett & Tsikoudi 1997).

The observed X-ray and EUV flares are gradual flares; we are

unlikely to have detected impulsive flares. We can assume, in

analogy with the Sun, that the observed X-ray and EUV emission

is of thermal origin. EXOSAT observations of X-ray flares

(Pallavicini et al. 1990) on late-type stars have shown both

gradual and impulsive flares (duration less than a few minutes).

The X-ray flares detected by Schmitt (1994), using the ROSAT

survey data were also all long-duration flares. He reports X-ray

flares on a variety of main-sequence stars, like 36 Dra, UV Cet

(Gl 65), EV Lac (Gl 873), etc. However, the short impulsive-type

events seen in EXOSAT data are extremely unlikely to be seen in

the ROSAT survey data. This is because the survey data is taken in

scanning mode, which results in, for the case of the PSPC, about

30 s of data exposure each scan period (about 96 min). That is,

there is an exposure `efficiency' of only about 0.5 per cent.

Therefore, we are much more `sensitive' to longer duration events,

particularly events that last longer than a single ROSAT orbital

period (,96 min).

The X-ray and the EUV peak flare luminosities are tightly

correlated (Fig. 6), a result which indicates that flares in our

sample are largely in the 2 million K and hotter temperature range.

The simultaneous flares correlate rather tightly whereas the non-

simultaneous ones show some slight increase in scatter. The fact

that these non-simultaneous flares are not greatly different in ratio

from the simultaneous flares could support the association

between the offset EUV flare and the XRT events. EUV-only

flares would be much cooler (.1 million K) and would be

unlikely to produce the `correct' EUV-to-X-ray flux ratio seen

Figure 7. For the 30 X-ray flares and the 14 quiescent measurements we show how the ratio of X-ray-to EUV emission varies with the X-ray luminosity.

1144 V. Tsikoudi et al.

q 2000 RAS, MNRAS 319, 1136±1146

here. We found no correlation between X-ray flare luminosity and

stellar colours.

Our results show that the non-flaring (quiescent) X-ray and

EUV luminosities are strongly correlated, for dKe and dMe stars,

over three orders of magnitude in both X-ray and EUV [Fig. 4(a)].

This tight correlation between simultaneous X-ray and EUV

observations indicates that there are no significant temperature

differences or gradients across our sample. However, we will

mention below that the exact ratio of X-ray to EUV emission is not

completely in agreement with the WFC-only filter ratio.

The portion of the LBol which is emitted in the X-rays ranges

between 2:6 � 1024 and 1:7 � 1023 (Table 2), for stars of spectral

type dK0±dM6. The ratio LX to LBol is considered a strong activity

indicator. For the Sun this ratio is about 1025; for later-type stars,

K±M, it becomes 1024±1023. For the highly active non-main-

sequence stars like the RS CVn's LX/LBol is up to about 1023. The

maximum value of LX to LBol is ,1023 and is claimed to be the

saturated coronal heating rate (Haisch & Schmitt 1994).

We have also estimated the hardness ratio, XRT/EUV, for each

flare and we compared it to the X-ray flare luminosity. In Fig. 7

we show all the X-ray flares which had an EUV counterpart. Also

shown in the figure are the quiescent X-ray-to EUV ratios. This

plot shows a fair degree of scatter, but does not show any

significant trend with increasing X-ray luminosity. However, the

ratio of X-ray to EUV of 3±10 when inverted and placed on

Fig. 1(b) (0.1±0.33) suggests a two-component model with the hot

component around 2±10 � 106 K ± a result that is somewhat at

odds with our earlier estimate based just on the WFC filter ratios.

This might suggest that the XRT1PSPC is `missing' some flux

from the hotter temperature component (we need to reduce the

ratio to about 0.03±0.06 for complete agreement). Two flares do

stand out from the rest, the second flares on Gl 644 and SAO

86592, which both produce rather weak EUV responses (and both

are non-simultaneous), but the very large error-bars on these

points (caused by the very small EUV flux) means that these flares

cannot be considered to be different from the rest. This plot, which

is essentially equivalent to the earlier plots in Figs 4(a) and 6,

seems to suggest that the quiescent X-ray and EUV emission

mechanism and the X-ray and EUV flare emission are following

the same relationship. Since the 30 X-ray/EUV flares cover

essentially the same luminosity range as the quiescent emission,

and since they are obviously flares, the apparent similarity of the

X-ray-to EUV ratio between flares and quiescent emission could

be indicating that the quiescent emission is just the integrated

emission from a series of unresolved underlying flares. It is

perhaps interesting to note that the EUV/X-ray filter ratios for the

flares do fall into two temperature bands which broadly agree with

the two-temperature spectral fits for similar active stars (Schmitt

et al. 1990).

The ROSAT survey observations offer the opportunity to

investigate the coronal X-ray variability of stars during their

non-flaring state. The time-interval of the survey observations

over any source is long enough (typically 2±3 d), to statistically

establish significant emission changes because it includes 30±45

telescope passes over the source. It is, therefore, possible to

observe how the coronal emission changes from one scan to the

next. The variability analysis which we applied to the ROSAT

survey observations, detected quiescent EUV variability on nine

of the 14 sample stars. Seven of the EUV-variables are also X-ray

variables in the ROSAT±PSPC survey. This is an important result

because for half of the sample stars we find `simultaneous' X-ray

and EUV short-term variability. It indicates that the coronal

quiescent variability is real and implies that the stellar coronae are

in a frequent state of low-level, short-term, activity. Again, this

could be revealing that the coronal quiescent emission is just the

integrated emission from a series of unresolved underlying flares.

For the flares themselves, we have recorded 31 X-ray flares in a

total of 36.33 survey days or 0.85 flares/day. This does include the

nearly 8-d long survey coverage of BY Dra (much longer than the

normal 2-d coverage because of its location near to the northern

Ecliptic pole). Taking out BY Dra and its four flares, we have 27

flares in 28.59 d (0.94 flares/day). These numbers convert to an

X-ray flare event, on average, every 1.06±1.17 survey days for our

14 stars. Given these numbers, the chances of detecting an

unrelated flare in the EUV (assuming the same flare rate as in the

X-ray) within ^3.2 h is about one in four. There are 15 EUV flares

seen, some not very significant, but given the statistical

uncertainty we still feel that the number of non-simultaneous

EUV flares is still significant. The most interesting case in our

sample is Gl 644 (Kellett & Tsikoudi 1999).

It appears that the stellar coronae of main-sequence stars are

X-ray and EUV variable on different scales. Although coronal

changes on the micro-flaring-scale have not been established, as

yet (Pallavicini et al. 1990), changes on the scale of milliflaring

are detected, as the analysis of the ROSAT survey observations of a

large number of stars has indicated (Kellett & Tsikoudi 1997).

Similar phenomena have been observed on the Sun. Yohkoh

observations have established that the solar X-ray corona is never

static (Acton, Feldman & Brunner 1992). The solar corona is

observed to be in a constant state of brightening and fading which

implies that there is constant small-scale flaring.

Both the X-ray and EUV observational results indicate that the

stellar coronae along the main sequence are bright and highly

dynamic. The coronal X-ray and EUV quiescent variability implies

that there is a frequent occurrence of small flare events (some-

times unresolved) in the stellar coronae. These results provide

evidence that the stellar coronae could be continuously heated by

such small-scale flaring.

AC K N OW L E D G M E N T S

One of us (VT) wishes to thank the staff of the Max-Planck-

Institut fuÈr Extraterrestrische Physik for their help when visiting

and for granting us access to the RASS data archive. We would

also like to thank the referee (Dr J. Pye) for his many helpful and

constructive comments which led to considerable improvements in

the final version of this paper.

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