A&A 389, 228-238 (2002)
DOI: 10.1051/0004-6361:20020529

High-resolution X-ray spectroscopy of Procyon by Chandra and XMM-Newton

A. J. J. Raassen1,2 - R. Mewe1 - M. Audard3 - M. Güdel3 - E. Behar4 - J. S. Kaastra1 - R. L. J. van der Meer1 - C. R. Foley5 - J.-U. Ness6


1 - SRON National Institute for Space Research, Sorbonnelaan 2, 3584 CA Utrecht, The Netherlands
2 - Astronomical Institute "Anton Pannekoek'', Kruislaan 403, 1098 SJ Amsterdam, The Netherlands
3 - Paul Scherrer Institut, Würenlingen & Villigen, 5232 Villigen PSI, Switzerland
4 - Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA
5 - Mullard Space Science Laboratory, University College London, Surrey, RH5 6NT, UK
6 - Universität Hamburg, Gojenbergsweg 112, 21029 Hamburg, Germany

Received 27 August 2001 / Accepted 5 April 2002

Abstract
We report the analysis of the high-resolution soft X-ray spectrum of the nearby F-type star Procyon in the wavelength range from 5 to 175 Å obtained with the Low Energy Transmission Grating Spectrometer (LETGS) on board Chandra and with the Reflection Grating Spectrometers (RGS) and the EPIC-MOS CCD spectrometers on board XMM-Newton. Line fluxes have been measured separately for the RGS and LETGS. Spectra have been fitted globally to obtain self-consistent temperatures, emission measures, and abundances. The total volume emission measure is ${\sim} 4.1 \times 10^{50}$ cm-3 with a peak between 1 and 3 MK. No indications for a dominant hot component ($T \ga 4$ MK) were found. We present additional evidence for the lack of a solar-type FIP-effect, confirming earlier EUVE results.

Key words: stars: individual: Procyon, $\alpha$ Canis Minoris - stars: coronae - stars: late-type - stars: activity - X-rays: stars


1 Introduction

Magnetized hot outer atmospheres (coronae) are ubiquitous in cool stars (spectral classes F-M). The particular example of the solar corona has revealed rich details on coronal structures, thermal stratification, abundance patterns, and the physics of heating and mass motion. Nevertheless, in many other stars coronal phenomena not common to the Sun are regularly observed, such as persistent high-density ( $n_{\rm e}>10^{10}$ cm-3) coronal plasmas, persistent very hot gas (T>10 MK), or abundances at variance with solar values (Bowyer et al. 2000). The Sun's relatively modest magnetic activity is representative for a particular evolutionary state of a $1~M_{\odot}$star (Güdel et al. 1997), while most stellar objects regularly observed by X-ray satellites belong either to the more abundant low-mass classes with some exceptional magnetic activity (M dwarfs and young K dwarfs) or to tidally coupled binary systems with strongly enhanced magnetic dynamos (RS CVn or Algol-type binaries).
  \begin{figure} \hbox{ \psfig{figure=H3110F1a.ps,angle=-90,height=5.9truecm,width... ...gure=H3110F1c.ps,angle=-90,height=5.9truecm,width=18truecm,clip=} } \end{figure} Figure 1: From top to bottom the spectra of Procyon observed by RGS1, RGS2 and LETGS in the wavelength region from 10 to 37 Å.
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  \begin{figure} \par\psfig{figure=H3110F2a.ps,angle=-90,height=5.96truecm,width=1... ...figure=H3110F2c.ps,angle=-90,height=5.96truecm,width=18truecm,clip=}\end{figure} Figure 2: The Procyon spectrum observed by LETGS in the region 36-174 Å.
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High-resolution X-ray spectroscopy of such stellar systems now available from Chandra and XMM-Newton (Brinkman et al. 2000, 2001) has revealed coronal features clearly at variance with solar phenomena. However, to translate solar knowledge to stellar environments, it is important to study stars that are relatively similar to the Sun. Given the low X-ray luminosity of such stars, there are only a few in the solar neighborhood accessible to high-resolution X-ray spectroscopy. We present here a detailed analysis of the X-ray spectrum of Procyon, a nearby bright X-ray source with a coronal plasma of about 1-3 MK, exhibiting a cooler X-ray spectrum than magnetically active stars that have predominantly been studied so far with XMM-Newton and Chandra (Audard et al. 2001a,b; Güdel et al. 2001a,b; Mewe et al. 2001). The late-type (F5 IV-V) optically bright ( $m_{\rm v}=0.34$) star Procyon (with a faint white dwarf companion) at a distance of 3.5 pc has a line-rich coronal spectrum in the X-ray region. The mass of Procyon is $1.75~M_\odot$ and its radius $2.1~R_\odot$ (Irwin et al. 1992). The high-resolution spectrum of Procyon has been studied earlier using EUVE (Drake et al. 1995; Schrijver et al. 1995; Schmitt et al. 1996) and by means of the LETGS on board Chandra (Ness et al. 2001). Ness et al. focussed their efforts on the density-sensitive and temperature-sensitive lines of C V, N VI, and O VII only. Here we present an extended investigation of the LETGS spectrum covering the total spectral range from 5-175 Å together with the analysis of the RGS spectra from 5-37 Å. In Sect. 2 we describe the observations. Section 3 (Analysis) is divided into a part on global fitting (3.2) based on the total spectrum and a part that contains consistency checks based on individual line measurements (3.3).


 

 
Table 1: Wavelengths and fluxes for RGS1, RGS2, and LETGS, together with the line identifications from MEKAL and KELLY.
RGS1 RGS2 LETGS Line identificationsa
$\lambda $(Å) fluxb $\lambda $(Å) fluxb $\lambda $(Å) fluxb MEKAL fluxc KELLY Ion
gap - 13.454(9) 0.17(4) 13.450(13) 0.14(5) 13.448 0.14 13.447 Ne IX
gap - 13.701(6) 0.17(4) 13.690(15) 0.12(5) 13.700 0.10 13.700 Ne IX
15.018(7) 0.20(5) 15.024(9) 0.26(4) 15.015(18) 0.21(6) 15.014 0.16 15.013 Fe XVIId
15.176(13) 0.12(3) 15.161(21) 0.08(3) - - 15.175 0.03 15.176 O VIII
- - - - 15.207(28) 0.13(5) - 0.09 - Fe XVIIe
15.281(12) 0.11(3) 15.258(12) 0.09(5) - - 15.265 0.04 15.260 Fe XVII
16.008(3) 0.17(3) 16.013(9) 0.20(5) 16.008(9) 0.18(5) 16.003 0.15 16.007 O VIIIf
16.776(9) 0.16(3) 16.776(15) 0.12(4) 16.790(14) 0.13(5) 16.780 0.10 16.775 Fe XVII
17.047(9) 0.23(4) 17.044(14) 0.19(5) 17.054(12) 0.22(9) 17.055 0.12 17.051 Fe XVII
- - 17.099(12) 0.25(5) 17.102(9) 0.29(9) 17.100 0.10 17.100 Fe XVII
17.402(12) 0.07(3) 17.402(28) 0.06(3) 17.396(17) 0.08(5) 17.380 0.04 17.396 O VII
17.765(9) 0.10(3) 17.772(9) 0.11(4) 17.769(13) 0.14(5) 17.770 0.12 17.768 O VII
18.624(4) 0.30(4) 18.637(6) 0.37(5) 18.629(5) 0.39(8) 18.628 0.39 18.627 O VII
18.970(2) 1.61(8) 18.975(2) 1.78(11) 18.972(2) 1.83(15) 18.973 1.93 18.969 O VIII
20.918(27) 0.04(3) gap - 20.905(22) 0.14(7) 20.910 0.06 20.910 N VII
21.596(2) 2.36(11) gap - 21.597(2) 3.01(25) 21.602 3.35 21.602 O VII
21.797(4) 0.53(6) gap - 21.792(5) 0.90(14) 21.800 0.80 21.804 O VII
22.098(2) 2.20(11) gap - 22.089(2) 2.57(23) 22.100 2.33 22.101 O VII
- - 24.780(3) 0.88(7) 24.790(4) 0.80(14) 24.781 0.76 24.781 N VII
- - 24.907(22) 0.08(4) 24.906(11) 0.18(9) 24.900 0.08 24.898 N VI
27.001(11) 0.16(4) 26.994(12) 0.14(4) 26.979(19) 0.18(9) 26.990 0.08 26.990 C VI
28.460(7) 0.30(5) 28.465(6) 0.39(6) 28.470(6) 0.49(12) 28.470 0.39 28.466 C VI
28.785(4) 0.69(9) 28.775(6) 0.71(8) 28.785(5) 0.88(15) 28.790 0.77 28.787 N VI
29.078(9) 0.24(5) 29.084(9) 0.29(7) 29.082(12) 0.29(10) 29.090 0.32 29.084 N VI
29.524(6) 0.43(6) 29.520(8) 0.38(6) 29.546(11) 0.43(13) 29.530 0.41 29.534 N VI
30.445(12) 0.20(5) 30.446(13) 0.22(5) 30.450(25) 0.18(11) 30.448 0.22 30.448 Ca XI
31.027(16) 0.18(5) 31.021(12) 0.15(5) 31.054(15) 0.22(10) 31.015 0.19 31.015 Si XII
- - 33.490(26) 0.15(7) 33.510(18) 0.17(10) - - - -
33.724(2) 3.49(17) 33.726(2) 4.15(30) 33.731(2) 4.02(32) 33.736 4.64 33.736 C VI
34.967(12) 0.19(6) 34.962(15) 0.17(7) 34.959(15) 0.27(17) 34.970 0.22 34.973 C V
35.198(30) 0.13(6) 35.193(12) 0.27(8) 35.188(18) 0.29(15) 35.212 0.12 35.212 Ca XI
35.566(16) 0.15(7) 35.562(12) 0.26(7) 35.566(16) 0.23(14) 35.576 0.27 35.576 Ca XI
35.682(8) 0.37(8) 35.676(9) 0.38(8) 35.672(9) 0.53(16) 35.665 0.28 35.665 S XIII
36.374(12) 0.35(10) 36.372(15) 0.28(7) 36.399(15) 0.34(14) 36.398 0.25 36.398 S XII
36.544(19) 0.15(6) 36.561(19) 0.29(9) 36.547(15) 0.24(13) 36.563 0.24 36.563 S XII
a Identifications for MEKAL (Mewe et al. 1995) and KELLY (Kelly 1987) identical.
b Observed flux in 10-4 photons/cm2/s.
c Model flux in 10-4 photons/cm2/s, obtained from 3-T fitting of LETGS (see Sect. 3.2).
d Fe XVII lines are strongly mixed with Fe XVI satellite lines.
e Line not split up; mixture of 15.175 Å and 15.265 Å from O VIII and Fe XVII, respectively.
f Small contamination by Fe XVIII possible.
Values in parentheses are statistical 1$\sigma $ uncertainties in the last digits.


2 Observations

The spectrum of Procyon was obtained during 70.4 ksec (on November 8, 1999) by the High Resolution Camera (HRC-S) and the LETGS on board Chandra. The HRC-S contains three flat detectors, each 10 cm long. LETGS consists of 180 grating modules. The LETGS spectrum covers the range from 5 to 175 Å. The LETGS spectra were summed over the +1 and -1 orders and contain also the higher orders. The higher orders are fitted in the model calculations, but can be neglected for Procyon. The curve of the effective area as a function of wavelength is complicated because of the presence of absorption edges (e.g., around 42 Å) and gaps between 62-65 Å and 52-56 Å (+1 and -1 order, respectively) due to the gaps between the detector plates. We use the SRON values (based on calibration by van der Meer et al. 2002), which agree within about 5-10% with the CXC values, as given in the Chandra LETGS Calibration Review of 31 Oct. 2001[*]. The wavelength resolution is $\Delta\lambda\sim0.06$ Å (FWHM). The wavelength uncertainty in the calibration is a few mÅ below 30 Å and about 0.020 Å in the region above 30 Å. The spectra are background subtracted. The statistical errors in the line fluxes include errors from the background. For further instrumental details see also Brinkman et al. (2000).

Later (on October 22, 2000), the spectrum of Procyon was observed by XMM-Newton using the RGS and EPIC-MOS. The total observing time was $\approx$107 ksec; however, due to large solar flare activity at the end of the observation, we removed 16.7 ksec of data, leaving a total of 90.5 ksec of "good'' data. In XMM-Newton three telescopes focus X-rays onto three EPIC cameras (two MOS and one pn). About half of the photons in the beams of two telescopes (Turner et al. 2001) are diffracted by sets of reflecting gratings and are then focussed onto the RGS detectors. The RGS spectral resolution is $\Delta\lambda\sim0.07$ Å, with a maximum effective area of about 140 cm2 around 15 Å. The wavelength uncertainty is 7-8 mÅ. The first spectral order has been selected by means of the energy resolution of the individual CCDs. For further details see den Herder et al. (2001).

The data were processed by the XMM-Newton SAS using the calibration of February 2001. The RGS cover the range from 5 to 37 Å. The EPIC spectra, which have a lower resolution but higher sensitivity, are used to constrain the high-temperature part of Procyon's EM distribution. Because of the high resolution of the grating spectrometers we will focus on the spectra from these instruments. In Fig. 1, we show the RGS spectra together with an extract of the LETGS spectrum covering the wavelength range from 10 to 37 Å. No notable features are observed below 10 Å in the LETGS and RGS spectra. However, the EPIC-MOS detects the H- and He-like lines of Mg. The remaining part of the LETGS spectrum is shown in Fig. 2. From Fig. 1 the gaps in the two RGS spectra due to CCD failure of CCD 7 (RGS1) and 4 (RGS2) are obvious.

3 Analysis

3.1 Detailed analysis of the spectra

The spectral lines from all three instruments have been measured individually. We folded monochromatic delta functions through the instrumental response matrices in order to derive the integrated line fluxes. No additional width was needed to fit the shape of the lines. A constant "background" level was adjusted in order to account for the real continuum and for the pseudo-continuum created by the overlap of several weak, neglected lines. In Table 1, we have collected the measured wavelengths and fluxes of the emission lines in the RGS instruments together with those in the LETGS in the similar wavelength range. The fluxes among the three data sets, as collected in Table 1, are in good agreement in view of the systematic uncertainties in the calibration. However for some lines deviations appear, which are caused by gaps between the individual CCDs. In this wavelength range (below 40 Å) the identification is in general straightforward. The dominating lines are strong and belong to H- and He-like ions for which atomic parameters are well known. Although XMM-Newton and Chandra observed Procyon at different dates no strong differences in flux (Table 1) are noticed, resulting in the conclusion that the coronal emission of Procyon did not vary strongly from one observation to the other.

Table 2 contains the same information as Table 1 for LETGS lines which occur above 37 Å. The extracted fluxes are as measured at Earth. Therefore they are not corrected for interstellar absorption which is of the order of 4% at 100 Å, 6% at 125 Å, 10% at 150 Å, and 15% at 175 Å. We added one Fe line (Fe X at 174.69 Å) that was observed in an offset observation of Procyon (obsID = 1224; 14.8 ksec). For that line the effective area was obtained by extrapolation. The line flux ratio of that line compared to the line at 171.075 Å in the offset observation was used to establish the flux value.

In both tables, we have compared the measured wavelengths with the wavelengths in various atomic databases: the MEKAL (Mewe et al. 1985, 1995) code[*], KELLY (Kelly 1987) and the database of the National Institute of Standards and Technology (NIST), which is also available on the web[*]. We have also compared with a list of lines observed in the solar corona (Doschek & Cowan 1984, hereafter D&C). Further we compare our measured iron lines with the results from laboratory experiments such as the Lawrence Livermore National Laboratory's Electron Beam Ion Trap (EBIT) (see Beiersdorfer et al. 1999 and Lepson et al. 2002 for Fe VIII-X and Lepson et al. 2000 for Fe XII-XIII). A number of lines in Table 2 (see note "d") are in close wavelength agreement to lines identified in EBIT. Finally in Table 1 the fluxes, from the multi-temperature global fitting of Sect. 3.2, have been added.

Some possible line identifications have been omitted from Table 2, due to the absence of comparable lines belonging to the same multiplet or ion (Table 3) or due to ambiguity of the identification of lines in atomic databases (Kelly 1987). The latter concerns lines at 60.989 Å (Si VII, VIII, & IX) and 61.852 Å (Si VIII & IX).

Earlier benchmarks with a solar flare spectrum (Phillips et al. 1999) and with RGS and LETGS spectra of Capella (Audard et al. 2001a; Mewe et al. 2001) have already shown that the current atomic databases are lacking quite a number of spectral lines for L-shell transitions of Ne, Mg, Si, and S, that appear in the long-wavelength region above about 40 Å. This is illustrated by the many identifications present in the third column (KELLY), which are absent in MEKAL. For the Fe L-shell Behar et al. (2001) have shown that the HULLAC atomic data are fairly accurate.

 

 
Table 2: Wavelengths and fluxes for LETGS above 36.5 Å, together with the line identifications from MEKAL, KELLY, and D&C.
LETGS Line identificationsa
  MEKAL KELLY D&C
$\lambda $(Å) fluxb $\lambda $(Å) Ion $\lambda $(Å) Ion $\lambda $(Å) Ion
39.276 0.63(14) 39.300 S XI 39.300 S XI 39.30 S XI
        39.264 Si X    
        39.305 Si X    
40.263 2.29(36) 40.270 C V 40.268 C V 40.27 C V
40.718 1.88(42) 40.730 C V 40.731 C V 40.73 C V
41.475 1.07(29) 41.470 C V 41.472 C V 41.47 C V
    41.480 Ar IX 41.480 Ar IX    
42.543 1.29(28) 42.530 S X 42.543 S X 42.53 S X
42.810 0.33(17) -   42.826 Si XI -  
43.743 0.54(8) 43.740 Si XI 43.763 Si XI 43.74 Si XI
44.014 0.43(8) 44.020 Si XII 44.021 Si XII 44.02 Si XII
44.150 0.67(10) 44.165 Si XII 44.165 Si XII 44.17 Si XII
44.218 0.52(10) 44.249 Si IX 44.215 Si IX 44.22 Si IX
45.677 0.20(4) 45.680 Si XII 45.692 Si XII 45.68 Si XII
        45.684 Si X    
46.283 0.25(7) 46.300 Si XI 46.300 Si XI 46.30 Si XI
46.391 0.40(8) 46.410 Si XI 46.401 Si XI 46.41 Si XI
47.242 0.46(8) 47.280 Mg X 47.310 Mg X 47.31 Mg X
        47.231 Mg X    
        47.249 S IX 47.25 S IX
47.452 0.48(8) 47.500 S IX 47.433 S IX 47.43 S IX
        47.518 S IX    
        47.453 Si XI    
47.642 0.49(8) 47.654 S X 47.655 S X 47.65 S X
        47.653 Si XI    
47.774 0.34(7) 47.793 S X 47.791 S X 47.79 S X
47.883 0.24(8) 47.896 Mg X 47.905 S X 47.90 Mg X
        47.899 Si XI    
48.720 0.23(6) 48.730 Ar IX 48.73 Ar IX 48.73 Ar IX
49.109 0.33(8) -   49.119 S IX 49.12 S IX
49.207 1.44(14) 49.220 Si XI 49.222 Si XI 49.22 Si XI
    49.180 Ar IX 49.18 Ar IX 49.18 Ar IX
49.324 0.32(7) -   49.328 S IX -  
49.696 0.29(7) -   49.701 Si X -  
49.975 0.28(6) 50.019 Si VIIIc 50.019 Si VIII -  
50.327 0.51(8) -   50.333 Si X -  
50.520 1.68(15) 50.530 Si X 50.524 Si X 50.53 Si X
50.686 1.30(14) 50.690 Si X 50.691 Si X 50.69 Si X
52.306 0.75(11) 52.300 Si XI 52.296 Si XI 52.30 Si XI
52.594 0.35(8) 52.615 Ni XVIII 52.615 Ni XVIII -  
        52.611 Si IX    
52.715 0.30(7) 52.720 Ni XVIII 52.720 Ni XVIII 52.70 S VIII
52.772 0.30(7) -   52.756 S VIII -  
        52.789 S VIII    
52.898 0.30(7) 52.911 Fe XV 52.911 Fe XV 52.87 Fe XV
54.133 0.56(13) 54.118 S VIII 54.118 S VIII 54.12 S VIII
    54.142 Fe XVI 54.142 Fe XVI 54.15 Fe XVI
54.180 0.31(10) 54.180 S IX 54.175 S IX 54.18 S IX
54.546 0.54(13) -   54.571 Si X -  
        54.566 S VIII    
54.700 0.45(11) 54.728 Fe XVI 54.728 Fe XVI 54.70 Fe XVI
55.094 0.68(15) 55.060 Mg IX 55.060 Mg IX 55.06 Mg IX
        55.094 Si IX 55.12 Si IX
        55.116 Si IX    
        55.096 Si X    
55.270 0.88(25) 55.272 Si IX 55.272 Si IX 55.27 Si IX
        55.305 Si IX 55.31 Si IX
55.359 2.14(27) 55.356 Si IX 55.356 Si IX 55.36 Si IX
        55.401 Si IX 55.40 Si IX
56.037 0.19(11) 56.000 Ni XIII 56.027 Si IX 56.03 Si IX
        56.081 S IX 56.08 S IX
56.836 0.20(8) -   56.804 Si X -  
57.741 0.80(35) -   57.736 Mg VIII -  
        57.778 Si IX    
57.856 0.78(35) 57.880 Mg X 57.876 Mg X 57.88 Mg X
    57.920 Mg X 57.920 Mg X 57.92 Mg X
61.020 1.41(25) 61.050 Si VIII 61.019 Si VIII 61.01 Si VIII
        61.038 Mg IX    
61.087 1.38(24) -   61.070 Si VIII 61.08 Si VIII
        61.088 Mg IX    
61.578 0.52(17) 61.600 S VIII 61.600 S VIII 61.60 S VIII
61.668 0.49(15) -   61.649 Si IX 61.66 Si IX
61.843 0.67(11) 61.841 Si IX 61.852 Si IX 61.84 Si IX
61.916 0.55(17) 61.912 Si VIII 61.914 Si VIII 61.91 Si VIII
        61.895 Si VIII 61.90 Si VIII
62.748 0.53(17) 62.755 Mg IX 62.751 Mg IX 62.76 Mg IX
    62.699 Fe XIIId 62.694 Fe XIII    
62.849 0.38(11) 62.879 Fe XVI 62.879 Fe XVI 62.88 Fe XVI
    62.800 Fe Xd 62.8 Fe X    
63.161 0.64(13) 63.153 Mg X 63.152 Mg X 63.15 Mg X
63.283 0.94(15) 63.294 Mg X 63.295 Mg X 63.29 Mg X
63.390 0.38(8) 63.314 Mg X 63.304 S VIII 63.40 Mg VII
        63.396 Mg VII    
63.720 0.58(11) 63.719 Fe XVI 63.719 Fe XVI 63.71 Fe XVI
        63.732 Si VIII 63.73 Si VIII
63.921 0.39(10) -   -   -  
64.135 0.44(11) -   -   -  
65.677 0.38(11) 65.672 Mg X 65.672 Mg X 65.67 Mg X
65.826 0.49(13) 65.840 Mg X 65.847 Mg X 65.84 Mg X
        65.822 Ne VIII    
65.884 0.41(10) 65.905 Fe XIId 65.905 Fe XII -  
        65.892 Ne VIII -  
66.057 0.28(10) 66.047 Fe XIId 66.047 Fe XII 66.04 Fe XII
66.255 0.64(14) -   66.259 Ne VIII -  
66.352 0.63(13) -   66.330 Ne VIII -  
67.161 0.48(11) 67.132 Mg IX 67.135 Mg IX 67.13 Mg IX
67.255 0.87(18) 67.233 Mg IX 67.239 Mg IX 67.22 Mg IX
        67.291 Fe XIId    
67.375 0.68(14) 67.350 Ne VIII 67.382 Ne VIII 67.35 Ne VIII
69.646 2.03(21) 69.658 Si VIII 69.632 Si VIII 69.66 Si VIII
    69.660 Fe XV 69.66 Fe XV    
69.827 1.05(14) 69.825 Si VIII 69.790 Si VIII 69.83 Si VIII
70.046 0.70(11) 70.020 Si VII 70.027 Si VII 70.03 Si VII
    70.010 Fe XII 70.01 Fe XII    
    70.054 Fe XV 70.054 Fe XV 70.05 Fe XV
71.929 0.69(13 -   71.901 Mg IX 71.92 Mg IX
        71.955 Si VII    
72.034 0.43(13) 72.030 Mg IX 72.027 Mg IX 72.03 Mg IX
72.302 1.44(18) 72.311 Mg IX 72.312 Mg IX 72.31 Mg IX
    72.310 Fe XI 72.310 Fe XI    
72.668 0.73(14) 72.663 S VII 72.663 S VII 72.66 S VII
72.871 0.58(15) 72.850 Fe IXd 72.850 Fe IX -  
73.478 0.47(11) -   73.470 Ne VIII -  
    73.471 Fe XV 73.471 Fe XV 73.47 Fe XV
73.555 0.43(11) 73.560 Ne VIII 73.563 Ne VIII 73.56 Ne VIII
74.860 1.10(18) 74.854 Mg VIII 74.858 Mg VIII 74.85 Mg VIII
    74.845 Fe XIIId 74.845 Fe XIII    
75.035 1.05(18) 75.034 Mg VIII 75.034 Mg VIII 75.03 Mg VIII
75.978 0.47(11) 76.006 Fe Xd 76.006 Fe X 76.02 Fe X
76.038 0.77(13) -   -   -  
76.507 0.32(11) 76.502 Fe XVI 76.502 Fe XVI 76.51 Fe XVI
76.862 0.55(13) -   -   76.87 Fe XVI
77.740 1.11(18) 77.741 Mg IX 77.737 Mg IX 77.74 Mg IX
78.733 0.71(14) 78.717 Ni XI 78.744 Ni XI 78.72 Ni XI
    78.769 Fe Xd 78.769 Fe X    
79.483 0.58(13) 79.488 Fe XII 79.488 Fe XII 79.49 Fe XII
80.017 0.38(11) 80.022 Fe XII 80.022 Fe XII 80.02 Fe XII
80.236 0.54(14) -   80.255 Mg VIII -  
80.507 0.74(14) 80.501 Si VI 80.501 Si VI 80.50 Si VI
    80.510 Fe XII 80.510 Fe XII 80.51 Fe XII
80.751 0.54(14) -   80.725 Si VI -  
81.865 0.42(11) -   81.895 Si VII -  
82.420 0.48(11) 82.430 Fe IXd 82.430 Fe IX 82.43 Fe IX
82.667 0.97(17) 82.744 Fe XII 82.598 Mg VIII -  
82.808 0.35(10) 82.837 Fe XII 82.837 Fe XII -  
        82.822 Mg VIII    
83.337 0.38(11) -   83.358 Si VI -  
83.600 0.55(15) -   83.587 Mg VII 83.59 Mg VII
        83.611 Si VI    
83.764 0.47(17) 83.766 Mg VII 83.766 Mg VII 83.77 Mg VII
83.935 0.46(11) 83.959 Mg VII 83.959 Mg VII 83.96 Mg VII
        83.910 Mg VII 83.91 Mg VII
84.032 0.40(11) -   84.025 Mg VII 84.02 Mg VII
84.292 0.40(11) 84.292 Ne VII 84.292 Ne VII 84.29 Ne VII
        84.212 Ne VII    
84.433 0.39(11) -   -   -  
85.448 0.38(11) -   85.477 Fe XII 85.47 Fe XII
        85.407 Mg VII 85.41 Mg VII
86.765 1.13(17) 86.772 Fe XI 86.772 Fe XI 86.77 Fe XI
86.876 0.55(17) -   86.847 Mg VIII -  
87.021 0.46(14) 87.025 Fe XI 87.025 Fe XI 87.02 Fe XI
        87.017 Mg VIII    
88.087 1.68(20) 88.092 Ne VIII 88.092 Ne VIII 88.08 Ne VIII
88.893 0.68(14) -   -   -  
88.955 0.75(15) -   88.952 Mg VI -  
89.156 0.43(13) 89.185 Fe XI 89.185 Fe XI 89.18 Fe XI
90.719 0.59(13) -   90.708 Mg VII -  
90.989 0.43(10) 91.009 Fe XIV 91.009 Fe XIV -  
        90.955 Fe XVIIe    
91.529 0.52(13) 91.564 Ne VII 91.564 Ne VII 91.56 Ne VII
91.627 0.38(8) -   -   -  
91.777 0.58(13) 91.808 Ni X 91.790 Ni X 91.81 Ni X
92.155 0.51(13) -   92.123 Mg VIII -  
92.858 0.55(14) -   92.850 Ne VII -  
93.587 0.58(15) -   -   -  
94.001 1.70(24) 94.012 Fe Xd 94.012 Fe X 94.02 Fe X
95.339 0.90(18) 95.374 Fe X 95.374 Fe X 95.37 Fe X
    95.338 Fe Xd 95.338 Fe X -  
95.412 1.04(20) 95.483 Mg VI 95.483 Mg VI 95.48 Mg VI
        95.421 Mg VI 95.42 Mg VI
95.997 1.46(20) -   96.022 Si VI 96.02 Si VI
96.124 0.79(17) 96.122 Fe Xd 96.122 Fe X 96.12 Fe X
96.804 0.71(18) 96.788 Fe Xd 96.788 Fe X -  
97.104 0.34(15) 97.122 Fe Xd 97.122 Fe X 97.12 Fe X
97.486 0.78(17) 97.502 Ne VII 97.502 Ne VII 97.50 Ne VII
98.091 1.58(25) 98.115 Ne VIII 98.115 Ne VIII 98.13 Ne VIII
98.251 2.89(34) 98.260 Ne VIII 98.260 Ne VIII 98.26 Ne VIII
100.57 1.05(24) -   100.597 Mg VIII -  
102.85 0.90(22) 102.91 Ne VIII 102.911 Ne VIII 102.9 Ne VIII
103.07 1.68(27) 103.08 Ne VIII 103.085 Ne VIII 103.1 Ne VIII
103.54 2.08(32) 103.57 Fe IXd 103.566 Fe IX 103.6 Fe IX
103.88 0.72(22) -   -   -  
104.67 0.68(21) -   -   -  
104.78 0.86(21) 104.81 O VI 104.813 O VI -  
105.20 1.22(21) 105.21 Fe IXd 105.208 Fe IX 105.2 Fe IX
106.18 1.03(20) 106.19 Ne VII 106.192 Ne VII 106.2 Ne VII
111.23 1.49(28) -   111.198 Ca X -  
111.71 0.53(13) 111.57 Mg VI 111.552 Mg VI 111.6 Mg VI
    111.72 Mg VI 111.746 Mg VI 111.7 Mg VI
113.33 0.46(11) -   113.315 Fe VIII -  
113.77 0.60(20) -   113.763 Fe VIII -  
        113.793 Fe IX    
113.99 0.71(20) -   113.990 Mg V 114.0 Mg V
        114.029 Mg V    
114.54 0.48(14) -   114.564 Fe VIII -  
114.88 0.66(17) -   114.785 Mg V 114.8 Mg V
115.37 0.51(18) 115.33 Ne VII 115.33 Ne VII -  
    -   115.39 Ne VII -  
115.80 0.77(20) 115.83 O VI 115.826 O VI 115.8 O VI
115.89 0.46(17) -   -   -  
116.70 1.54(25) 116.69 Ne VII 116.693 Ne VII 116.7 Ne VII
116.87 0.54(14) -   -   -  
117.20 0.54(20) 117.20 Fe VIIId 117.197 Fe VIII -  
117.66 0.80(20) -   -   -  
119.31 0.46(13) -   -   -  
120.31 0.60(15) 120.33 O VII 120.331 O VII -  
122.49 0.73(18) 122.49 Ne VI 122.49 Ne VI 122.5 Ne VI
123.54 1.31(31) -   -   -  
124.51 0.98(25) -   -   -  
126.25 0.98(32) -   126.280 Mg V -  
127.53 0.98(25) -   -   -  
127.69 1.44(29) 127.66 Ne VII 127.663 Ne VII 127.7 Ne VII
129.86 1.07(34) 129.83 O VI 129.871 O VI 129.9 O VI
130.92 1.78(39) 130.94 Fe VIIId 130.941 Fe VIII 130.9 Fe VIII
131.21 1.78(35) 131.24 Fe VIIId 131.240 Fe VIII 131.2 Fe VIII
134.21 1.43(41) -   -   -  
135.48 0.97(35) 135.52 O V 135.523 O V 135.5 O V
136.78 1.91(53) -   -   -  
140.27 1.36(49) -   -   -  
141.04 0.79(22) 141.04 Ca XII 141.038 Ca XII 141.0 Ca XII
144.97 1.73(49) 144.99 Ni X 144.988 Ni X 145.0 Ni X
147.27 1.12(35) 147.27 Ca XII 147.278 Ca XII 147.3 Ca XII
148.36 11.3(10) 148.40 Ni XI 148.402 Ni XI 148.4 Ni XI
150.08 5.08(60) 150.10 O VI 150.1 O VI 150.1 O VI
151.52 2.36(39) -   151.548 O V -  
152.11 5.60(66) 152.15 Ni XII 152.153 Ni XII 152.2 Ni XII
154.14 2.73(42) 154.18 Ni XII 154.175 Ni XII 154.2 Ni XII
155.56 1.36(28) -   -   -  
156.14 1.66(32) -   156.140 Ne V -  
156.38 1.29(31) -   -   -  
157.68 2.76(42) 157.73 Ni XIII 157.730 Ni XIII 157.7 Ni XIII
158.33 2.08(35) 158.38 Ni X 158.377 Ni X 158.4 Ni X
158.78 1.47(34) -   158.770 Ni XIII -  
159.24 1.03(27) -   159.300 Si X 159.1 Ar XIII
159.58 1.83(43) -   -   -  
159.93 3.09(43) 159.94 Ni X 159.977 Ni X 159.9 Ni X
    159.97 Ni XIII 159.97 Ni XIII    
162.56 2.94(83) 162.56 N V 162.556 N V -  
164.11 3.53(74) 164.15 Ni XIII 164.146 Ni XIII 164.1 Ni XIII
167.43 3.9(10) 167.49 Fe VIII 167.486 Fe VIII 167.5 Fe VIII
167.59 3.9(12) 167.66 Fe VIII 167.656 Fe VIII -  
168.13 7.4(13) 168.17 Fe VIII 168.172 Fe VIII 168.2 Fe VIII
168.51 5.4(13) 168.54 Fe VIII 168.545 Fe VIII 168.5 Fe VIII
168.90 5.0(18) 168.93 Fe VIII 168.929 Fe VIII 168.9 Fe VIII
171.04 114(8) 171.08 Fe IX 171.075 Fe IX 171.1 Fe IX
174.49 118(24) 174.53 Fe X 174.53 Fe X 174.5 Fe X
a From MEKAL (Mewe et al. 1995), KELLY (1987), and D&C (solar line list of Doschek & Cowan 1984).
b Observed flux in 10-4 photons/cm2/s with in parentheses 1$\sigma $ uncertainty in the last digits.
c MEKAL placed it at 52.0 Å.
d Line identified in EBIT spectrum.
e From CHIANTI (Dere et al. 1997).


3.2 Global fitting and emission measure modeling

3.2.1 Multi-temperature fitting
We first characterize the thermal structure and the elemental composition of Procyon's corona. To this end, we fitted multi-T models using SPEX (Kaastra et al. 1996a) of the spectra (RGS+MOS and LETGS). For both the observations the calculations require two dominant temperature components. A third (small and not very significant) temperature component is needed to account for the lines of low stages of ionization, present in the LETGS spectrum. The reduced $\chi^2$ is relatively high (1.3-2) for the fits. This is due to a lack of lines in the MEKAL code and to the high resolution of the instrument. Small wavelength deviations (about 1-2 bins i.e. 0.02-0.04 Å) between lines in the spectrum and in the model are often present (see Table 2). This effect results in a sharp maximum and minimum in the value of the normalized difference between model and observation around the peak of the line (see also Fig. 4). The results of RGS and LETGS are very similar.

In Table 4 results for temperatures T (in MK), emission measures EM, and abundances are given. Statistic 1$\sigma $ uncertainties are given in parentheses. The emission measure is defined as $EM = n_{\rm e} n_{\rm H} V$, where V is the volume contributing to the emission and for solar abundances the hydrogen density $n_{\rm H} \simeq 0.85 n_{\rm e}$. The temperatures and emission measures of all spectra show a dominant region between 1 and 2.5 MK. The two dominant temperature components are about 1.2 and 2.3 MK. Using EUVE, Schmitt et al. (1996) derived a DEM with a peak temperature of 1.6 MK based on Fe-lines only. This is in satisfactory agreement with our results.

The total emission measures summed over all temperature components are about $4.6(.4) \times 10^{50}$ cm-3 for LETGS and $3.9(.3) \times 10^{50}$ cm-3 for RGS+MOS. These are similar to the total EM of $4.5 \times 10^{50}$ cm-3 found by Schmitt et al. (1996).

The determination of abundances is complicated by several factors. The many weak L-shell lines, which are absent in the atomic code (see difference between Col. "MEKAL'' and "KELLY'' of Table 2) can produce a "pseudo-continuum'' (see e.g. Fig. 2a between 42 and 58 Å), which bias the determination of the real but very weak continuum. Several fits to the LETGS spectrum were made: a) to the total spectrum, b) to the total spectrum with selected lines in the wavelength range from 40 to 100 Å, to limit the influence of the inaccuracy of atomic data of Ne-, Mg-, and Si- L-shell lines, and c) to a line spectrum with lines of Table 1 and lines with a statistical significance $\ga$4$\sigma $ in the wavelength range above 40 Å (see Table 5). During our investigations the absolute (relative to H) abundances turned out to be very sensitive to the selected group of elements introduced in the fit procedure. This is especially true for the elements Ar and Ca. For these reasons no consistent absolute values of the abundances could be obtained. However, abundance ratios turn out to be much more robust. Therefore the abundance values are normalized to oxygen, and are given relative to their solar photospheric values (Anders & Grevesse 1989), except for iron. For Fe we use log $A_{\rm Fe}$ is 7.51 (see Drake et al. 1995) instead of 7.67 (Anders & Grevesse 1989). Here log $A_{\rm Fe}$ is the logarithmic of the Fe-abundance relative to log $A_{\rm H}=12.0$. The abundances presented in Table 4 are derived assuming the same abundances for the three temperature components. These are averaged over the different fits, together with their least-squares-fit standard deviations (within parentheses).


  

 
Table 3: Possible line identifications left out of Table 2. Col. $\lambda $: observed wavelengths from Table 2. Columns 2 and 3 give a possible identification which has not been given in Table 2, due to the absence of the lines, given in Cols. 4 and 5.
$\lambda $(Å) present ion missing ion
93.587 93.616 FeVIII 93.469 FeVIII
      108.077 FeVIII
98.583 98.548 FeVIII 98.371 FeVIII
103.88 103.937 FeXVIII 93.923 FeXVIII
103.88 103.904 MgV 110.859 MgV


We obtain abundance values between 0.7 and 2.4 relative to oxygen (e.g., some enhancement for Ne and Si). However, apart from statistical errors these values are also sensitive to systematic errors, due to changes in values of the solar photospheric abundances, where uncertainties of a factor of 2 cannot be excluded (e.g., Prieto et al. 2001; Grevesse & Sauval 1998). So we cannot obtain indications for a significant FIP effect (enhancement of elements with a low First Ionization Potential) as found for the solar corona (e.g., Feldman et al. 1992). This confirms the conclusions by Drake et al. (1995), based on relative abundances from EUVE observations. The abundances of C and N, relative to O are somewhat higher than the values obtained in the solar photosphere (Anders & Grevesse 1989). In the EUVE observations by Drake et al. (1995) no suitable C- and N-lines were present to constrain (relative) abundances.

Values for $n_{\rm e}$, given in Table 4, have been obtained by fitting to the O VII and N VI triplet lines. The C V lines have been omitted from this procedure because their intensities are sensitive for the stellar UV-radiative field, mimicing higher densities (Ness et al. 2001; Porquet et al. 2001).


 

 
Table 4: Best-fit parameters for a 3-T CIE model fit. Elemental abundances for the three instruments are given normalized to oxygen and relative to solar photospheric values given by Anders & Grevesse (1989), except for Fea. 1$\sigma $ uncertainties are given in brackets.
Parameter LETGS RGS+MOS
log $N_{\rm H}$ [cm-2] 18.06b 18.06b
T1 [MK] 0.63(.10) -
T2 [MK] 1.21(.07) 1.65(.15)
T3 [MK] 2.26(.12) 2.68(.22)
EM1 [1050 cm-3] 0.41(.14) -
EM2 [1050 cm-3] 2.45(.27) 3.0(.20)
EM3 [1050 cm-3] 1.72(.29) 0.9(.18)
$n_{\rm e}$2[1010 cm-3] 1.4 +1.5-0.6 1.5 +2.0-0.6
$n_{\rm e}$3[1010 cm-3] 0.2 +0.8-0.2 -
O/H 0.68(0.38) 0.76(0.33)
C/O 1.38(.24) 1.45(.29)
N/O 1.33(.10) 1.47(.5)
O/O 1.0 1.0
Ne/O 1.49(.21) 1.53(.27)
Mg/O 1.1(.5) -
Si/O 1.56(.36) -
S/O 0.69(.15) -
Fe/O 0.97(.31) 1.47(.22)
Ni/O 2.39(.27) -
a In logarithmic units, with log $_{10}~\rm H=12.00$; $\rm C=8.56$; $\rm N=8.05$; $\rm O=8.93$; $\rm Ne=8.09$; $\rm Mg=7.58$;
$\rm Si=7.55$; $\rm S=7.21$; $\rm Ar=6.56$; $\rm Ca=6.36$; $\rm Fe=7.51$ (see text); $\rm Ni=6.25$.
b See Linsky et al. (1995).


3.2.2 Temperature dependent emission measure modeling
To show the connectivity of the different temperature components we applied a differential emission measure (DEM) model of Procyon's corona using the various inversion techniques offered by SPEX (see Kaastra et al. 1996b). We applied the abundances obtained in Sect. 3.2.1. In Fig. 3 we give the results based on the regularisation method. Other inversion methods (smoothed clean, or polynomial) give statistically comparable results. The DEM modeling has been applied separately to RGS+MOS and to LETGS.

As a result we find a dominant emission measure of the order of 1050 cm-3 between 1-3 MK. The total emission measures are $3.5(.3) \times 10^{50}$ cm-3 for RGS+MOS and $4.5(.2) \times 10^{50}$ cm-3 for LETGS (in line with the multi-temperature fitting). Figure 3 allows us to conclude that there is no significant amount of EM at $T \ga$ 4 MK in the corona of Procyon. Schmitt et al. (1996) give an upper limit of 6 MK, based on EUVE observations. The EM observed at different times as well as lines fluxes in Table 1 show no significant variability.

Figure 4 shows fit residuals of parts of the LETGS spectrum fitted using this temperature-dependent emission measure modeling, i.e. applying the model of Fig. 3 (LETGS). From Fig. 4a we recognize large deviations in residual due to model insufficiencies and a pseudo-continuum (most fit residuals positive) due to the lack of weaker lines in current atomic databases in this wavelength range. Clear from Fig. 4b are the succeeding large positive and negative residuals around 148 and 171 Å, due to wavelength deviations of lines in the spectrum and the model.


  \begin{figure} \par\includegraphics[angle=-90,width=8.8cm,clip]{H3110F3.ps_online_couleur}\end{figure} Figure 3: EM ( $n_{\rm e} n_{\rm H} V$ in $10^{64}~\rm m^{-3}$) for RGS and LETGS (thick), using the regularisation algorithm. The relative abundances given in Table 4 have been applied.
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3.3 Consistency checks using individual lines

The question is whether the model insufficiencies influence our conclusions about temperatures, emission measures, and abundances as obtained in Sect. 3.2. Therefore we have also compared observed and model line fluxes. The advantage of this individual line approach is that we can select strong and unblended lines, for which the theoretical emissivities are quite well established.

For the short-wavelength region this is done for all lines (Table 1), while for the longer wavelength range only lines with a statistical significance $\ga$4$\sigma $ were used. For the latter the fluxes have been compared with the 3-T model as well as with the results from the DEM model. The values are given in Table 5. This table shows generally a good agreement between the observed flux and the 3-T flux and the flux from the DEM-modeling, summed over the T-bins. Most striking are the deviations for the Fe VIII lines around 131 and 168 Å. This is definitely due to a large deficiency in the atomic data used. From atomic physics grounds the line at 168.13 Å is the stronger, as observed in the spectrum and in laboratory experiments (Wang et al. 1984), but in our code this line turns out to be the weakest[*]. Another interesting feature is the contamination of the forbidden C V line - which is often used for density diagnostics - with Ar IX. Another clear example of blending is the line at 74.860 Å which contains Mg VIII and Fe XIII.


  \begin{figure} \par\includegraphics[height=8.8truecm,width=5.05cm,angle=-90,clip... ...egraphics[height=8.8truecm,width=5.05cm,angle=-90,clip]{H3110F4b.ps}\end{figure} Figure 4: Fit residuals ((observed - model)/error) of parts of the LETGS spectrum.
Open with DEXTER


 

 
Table 5: Observed line fluxes and fluxes obtained from the emissivity from the model.
LETGS Line identifications
Observed Model
$\lambda $(Å)   fluxa $\lambda $(Å) 3-T fluxb DEM-fluxb Ion
18.972 1.83(15) 18.973 1.93 1.80 N VII
21.597 3.01(25) 21.602 3.35 2.94 N VII
24.790 0.80(14) 24.781 0.76 0.70 N VII
33.731 4.02(32) 33.736 4.64 4.19 C VI
40.263 2.29(36) 40.270 2.03 1.40 C V
40.718 1.88(42) 40.730 1.33   C V
41.475 1.07(29) 41.470 0.68   C V
  41.480 0.28 0.31 Ar IX
43.743 0.54(8) 43.740 0.54 0.94 Si XI
44.150 0.67(10) 44.165 0.64 0.66 Si XII
47.242 0.46(8) 47.280 0.19 0.17 Mg X
47.452 0.48(8) 47.500 0.61 0.60 S IX
47.642 0.49(8) 47.654 0.20 0.31 S X
49.207 1.44(14) 49.220 0.43 0.95 Si XI
  49.180 0.41 0.37 Ar IX
50.520 1.68(15) 50.530 1.48 1.55 Si X
50.686 1.30(14) 50.690 1.50 1.58 Si X
52.306 0.75(11) 52.300 0.45 0.74 Si XI
61.020 1.41(25) 61.050 2.51 1.98 Si VIII
61.087 1.38(24)       Si VIIIc
63.283 0.94(15) 63.294 1.07 0.93 Mg X
69.646 2.03(21) 69.658 0.85 0.67 Si VIII
  69.660 1.14 1.11 Fe XV
74.860 1.10(18) 74.854 0.51 0.57 Mg VIII
  74.845 0.26 0.40 Fe XIII
75.035 1.05(18) 75.034 0.52 0.57 Mg VIII
77.740 1.11(18) 77.741 0.42 0.37 Mg IX
86.765 1.13(17) 86.772 0.71 0.45 Fe XI
88.087 1.68(20) 88.092 2.33 1.62 Ne VIII
98.251 2.89(34) 98.260 3.02 2.37 Ne VIII
105.20 1.22(21) 105.21 0.32 0.19 Fe IX
130.92 1.78(39) 130.94 0.29 0.20 Fe VIII
131.21 1.78(35) 131.24 0.41 0.29 Fe VIII
148.36 11.3(10) 148.40 11.0 15.5 Ni XI
150.08 5.08(60) 150.10 2.8 2.63 O VI
152.11 5.60(66) 152.15 2.9 6.0 Ni XII
167.43 3.9(10) 167.49 3.9 2.64 Fe VIII
167.59 3.9(12) 167.66 4.0 2.74 Fe VIII
168.13 7.4(13) 168.17 0.3 0.20 Fe VIII
168.51 5.4(13) 168.54 2.0 1.42 Fe VIII
168.90 5.0(18) 168.93 1.3 0.71 Fe VIII
171.04 114(8) 171.08 100 78 Fe IX
a Observed flux in 10-4 photons/cm2/s with in parentheses 1$\sigma $ uncertainty in the last digits.
b Model flux in 10-4 photons/cm2/s.
c Sum of two Si VIII lines to be compared with model flux.


We have measured line ratios of density-sensitive He-like triplets from the LETGS and RGS spectra, taking into account the photo-exciting UV flux (Porquet et al. 2001). Our results are consistent in both instruments ( $n_{\rm e} \approx 10^{10}$ cm-3) and similar to those of Ness et al. (2001) and our values given in Table 4. These results are also comparable to values obtained by Schrijver et al. (1995) and Schmitt et al. (1996) and to values for the Sun (Drake et al. 2000).

4 Conclusions

The RGS and LETGS spectra of the corona of Procyon below 40 Å are dominated by the H- and He-like transitions of C, N, and O and by Fe XVII lines. Above 40 Å the LETGS spectrum shows many L-shell lines of e.g., Ne, Mg, and Si, together with lines of Fe VIII-XIII of which the Fe IX line at 171.075 is very prominent. All methods applied in Sect. 3.2 to the spectra of the RGS+MOS and the LETGS show temperatures of the corona of Procyon between 1-3 MK. No indication for a considerably higher temperature component ($T \ga$ 4 MK) is found. The total EM obtained using RGS and LETGS is about $4.1(.5) \times 10^{50}$ cm-3. The EM distribution shows a smooth continuous structure without separated peak structures. Our results improve on those of Schmitt et al. (1996) who obtain an EM distribution with a maximum temperature around 1.6 MK and a cutoff beyond 6.3 MK. No significant variability of the coronal conditions took place between the observations by RGS and LETGS.

The abundances of C and N, relative to O are somewhat higher ($\sim$factor 1.5) than the values obtained in the solar photosphere (Anders & Grevesse 1989). The Fe abundance is about 1-1.5 $\times$ solar. No significance for a FIP effect, as observed in the solar corona (Feldman et al. 1992), is found. The same was concluded by Drake et al. (1995), based on EUVE observations. This result is an exception of the trends found by Audard et al. (2001c) for RS CVn systems and by Güdel et al. (2001c) for solar analogs. These authors have found indications for the evolution from an inverse FIP effect for highly active stars - via the absence of a FIP effect in intermediately active stars - towards a normal FIP effect for less active stars. Clearly, the weakly active star Procyon does not fit into this picture.

Acknowledgements

The Space Research Organization Netherlands (SRON) is supported financially by NWO. The PSI group acknowledges support from the Swiss National Science Foundation (grant 2100-049343). We are grateful to the calibration teams of the instruments on board XMM-Newton and Chandra. We thank Nancy Brickhouse and Jeremy Drake for their efforts to obtain a long LETGS observation. Finally, we are grateful to the referee for helpful comments.

References

 

Copyright ESO 2002