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Past and recent observations of the solar upper atmosphere at vacuum-ultraviolet wavelengths
Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
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Past and recent observations of the solar upper atmosphere at vacuum-ultraviolet wavelengths Klaus Wilhelm∗ Max-Planck-Institut fur Aeronomie (MPAE), Max-Planck-Strae 2, 37191 Katlenburg-Lindau, Germany Received 4 December 2001; received in revised form 25 June 2002; accepted 27 September 2002
Abstract Our understanding of the high-temperature solar atmosphere is to a large extent based on spectroscopic observations of emission lines and continuum radiation in the vacuum-ultraviolet (VUV) wavelength range of the electromagnetic spectrum. In addition, important contributions stem from soft X-ray measurements. The VUV radiation is produced by transitions of atoms and ions, or to some extent, of molecules. The atomic and ionic emission lines have formation temperatures between 10 000 K and several million kelvin, representative of the chromosphere, the transition region and the corona. The molecular lines and the continua originate in cooler regions of the Sun. Radiation at VUV wavelengths is strongly absorbed by the Earth’s atmosphere leading to important geophysical processes at high altitudes. In our context it means that this radiation can only be detected with instruments on sounding rockets and spacecraft above the atmosphere. Detailed studies of the spectral radiances together with atomic physics data furnish information on the electron density and temperature of the solar atmosphere, as well as on elemental abundances, whereas Doppler line-shift measurements show bulk plasma motions, turbulence, and ion temperatures. Highlights of the research in this
1. Introduction In this review the great signi
∗ Corresponding author. Tel.: +49-5556-979-423; fax: +495556-979-240. E-mail address: [email protected] (K. Wilhelm).
or spacecraft at altitudes above at least 150 km. This absorption, in turn, leads to important processes in the Earth’s ionosphere, including photoionization of N2 , O2 , NO, and O ? reat wavelengths shorter than 796, 1026, 1340, and 911 A, ? spectively, as well as photodissociation of N2 below 1270 A ? Consequently, a detailed knowledge and O2 below 2422 A. of the spectral irradiance of the solar VUV radiation is also essential for a quantitative treatment of the thermosphere. In this context, it is of special importance that the O2 absorp? exactly tion cross section has a deep minimum at 1216 A, at the wavelength of the strong and variable H I Ly line. This line is therefore absorbed at the 85 km level and does not signi
c 2002 Elsevier Science Ltd. All rights reserved. 1364-6826/03/$ - see front matter PII: S 1 3 6 4 - 6 8 2 6 ( 0 2 ) 0 0 2 8 5 - 7
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K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
Fig. 1. Spectral irradiance at 1 ua (astronomical unit) of the solar H I Ly pro
example might help to clarify this point. After the H I Ly line was
work accompanying the observational progress. As far as the observational aspects are concerned, there will be a bias towards the earliest and most recent results. There is just not enough space to present the complete development over more than
K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
the top of the photosphere, in perfect agreement with the value of 4430 K ±50 K found by Samain et al. (1975) from ? Most of the emission lines are the continuum around 1580 A. produced above this temperature minimum and stem from the chromosphere, transition region or corona of the solar atmosphere. The
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approximately v = 1 km s−1 , which are related to wavelengths by the Doppler formula v = c( − o )=o (where c is the speed of light and o is the rest wavelength of a speci
Spacecraftb
Rocket Rockets Rockets Solrad 1 Rocket Rocket OSO 1 Rocket Rocket OSO 3 OSO 3 OSO 4 OSO 4 Rocket OSO 6 OSO 7 Rocket Rocket Skylab/ATM Skylab/ATM Skylab/ATM Rocket OSO 8 OSO 8 Rockets Rockets Rockets SMM
Time
1955 1958/1959 1959 –1961 1960 1960 1961 1962 1963 1965 1967 1967 1967 1967 1969 1969 –1972 1971–1973 1973 1973 1973/1974 1973/1974 1973/1974 1974 1975 –1978 1975 –1978 1975 –1987 1979 –1982 1979/1980 1980 –1989
Telescope, grating (NI) Grating spectrometer (GI) Photoelectric detector (GI) Ly photometer Double dispersion (NI) GI with aluminium
Instrumentc 900 –3000 84 –1216 250 –1300 1050 –1350 500 –1550 170 –700 10 – 400 257–269 180 –2950 260 –1300 20 – 400 300 –1400 304, 1216 60 –385 285 –1385 170 – 400 200 –700 1200 –2100 280 –1340 171– 630 1150 –3940 584 1200 –2000 1025 –1216 1175 –1710 1216 –2200 605 – 633 1150 –3600
Wavelength ? ranged (A) 1 0.4 1.5 –3.0 Bandpass 0.1– 0.2 0.1– 0.2 0.86 — 0.4 2 0.6 1.6 — 0.04 3 0.8 0.02 0.4 – 0.5 1.6 0.027 0.05 15 0.02 0.02– 0.06 0.05 Bandpasses Pixel: (0.028) FWHM: 0.04, (0.02)
Spectral ? resolutione (A) Disk Disk Disk Disk Disk Disk Disk 60 4 Disk Disk 60 × 60 Disk Disk 35 × 35 10 × 20 20 7–10 Slit: 5 × 5 2–3 ... 10 2 × 60 Disk Slit: 2:5 × 3 2 0:5 × 0:8 61 Slit: 18 × 54 62
Angular resolutionf ( )
Jursa et al. (1955) Violett and Rense (1959) Hall et al. (1963) Kreplin et al. (1962) Detwiler et al. (1960) Austin et al. (1962) Behring (1970) Purcell et al. (1964) Burton et al. (1967) Hinteregger and Hall (1969) Neupert et al. (1968) Reeves and Parkinson (1970) Bowles et al. (1968) Behring et al. (1972) Huber et al. (1973) Underwood and Neupert (1974) Cushman and Rense (1976, 1977) Samain et al. (1975) Reeves et al. (1977) Tousey et al. (1977) Bartoe et al. (1977) Maloy et al. (1978) Bruner (1977) Bonnet et al. (1978) Brueckner and Bartoe (1983) Bonnet et al. (1980) Rottman et al. (1982) Woodgate et al. (1980)
Referenceg
Table 1 Selected solar VUV instrumentsa Hown on spacecraft (sub-orbital rockets, satellites, or space probes) and some of their typical performance characteristics
170 K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
Spacelab 2 Spacelab 2 Rocket Rockets Rockets Rockets Rockets Yohkoh UARS UARS Spartan 201 SOHO SOHO SOHO SOHO SOHO Rocket TRACE
CHASE (GI) HRTS (NI) Multi-layer telescopes EGS (NI) SERTS (GI) NIXT MSSTA SXT SUSIM (NI) SOLSTICE (NI) UCS (NI) SUMER (NI) CDS (GI,NI) EIT UVCS (NI) SEM (He II) XDT (Fe XIV) Multi-layer telescope
160 –1344 1176 –1700 44 –256 250 –1200 170 – 450 63.5 44 –1548 3– 45 1150 – 4100 1190 – 4200 1032, 1037, 1216 465 –1610 151–785 171–304 492–1287 304, 170 –700 205 –218 170 –1700
0.7 0.05 Bandpasses 1.7 FWHM: 0.05 – 0.08 FWHM: 1.4 Bandpasses Bandpass 10 1–2 Pixel: 0.25 Pixel: 0.044, (0.022) 0.3 Bandpasses 0.09 – 0.14 80, Bandpass ? red/blue 211:3 A: Bandpasses
15 1 1–1.5 Disk 7 0.75 ¡ 0:75 FWHM: 3 Disk Disk Corona: 30 × 150 1–2 4–6 2.5 –5 Corona: 7 Disk 5.2 Pixel: 0.5
Lang et al. (1990) Bartoe and Brueckner (1975) Walker et al. (1988) Woods and Rottman (1990) Neupert et al. (1992a) Golub et al. (1990) Hoover et al. (1991) Tsuneta et al. (1991) Brueckner et al. (1993) Rottman et al. (1993) Kohl et al. (1994) Wilhelm et al. (1995) Harrison et al. (1995) DelaboudiniSere et al. (1995) Kohl et al. (1995) Hovestadt et al. (1995) Sakao et al. (1999) Handy et al. (1999)
b Spacecraft
soft X-ray instruments are included. and spacecraft subsystem abbreviations: ATM, Apollo Telescope Mount; OSO, Orbiting Solar Observatory; SMM, Solar Maximum Mission; SOHO, SOlar and Heliospheric Observatory; TRACE, Transition Region and Coronal Explorer; UARS, Upper Atmosphere Research Satellite. c Instrument abbreviations: GI, Grazing incidence; NI, Normal incidence; CDS, Coronal Diagnostic Spectrometer; CHASE, Coronal Helium Abundance Spacelab Experiment; EGS, EUV Grating Spectrograph; EIT, Extreme-ultraviolet Imaging Telescope; HRTS, High Resolution Telescope and Spectrometer; MCS, MultiChannel Spectrometer; MSSTA, Multi-Spectral Solar Telescope Array; NIXT, Normal Incidence X-ray Telescope; SEM, Solar EUV Monitor; SERTS, Solar EUV Research Telescope and Spectrograph; SOLSTICE, SOLar-STellar Irradiance Comparison Experiment; SUMER, Solar Ultraviolet Measurements of Emitted Radiation; SUSIM, Solar Ultraviolet Spectral Irradiance Monitor; SXT, Soft X-ray Telescope; TRC, Transition Region Camera; UCS, Ultraviolet Coronal Spectrometer; UVCS, UltraViolet Coronagraph Spectrometer; UVS, UltraViolet Spectrometer; UVSP, UltraViolet Spectrometer and Polarimeter; XDT, XUV Doppler Telescope. d In some cases, not the full range is actually covered. e Spectral resolution and resolution elements are not distinguished systematically. Second-order values in parentheses. FWHM: Full-width at half-maximum. f Spatial resolution, resolution elements, or slit dimensions. g References with instrumental details.
a Some
1985 1985 1987 1988–1994 1989 – 1989 – 1991–1994 1991–2002 1991–2001 1991–2001 1993– 1995 – 1995 – 1995 – 1995 – 1995 – 1998 1998– K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189 171
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Fig. 2. Spectral radiance of a quiet-Sun region near the centre of the disk on 12 August 1996. The SUMER grating spectrometer superimposes spectra of the
2. Emission line spectra 2.1. Spectral line identi>cations Johann Wilhelm Ritter observed in 1801 the decomposition of silver chloride on the short-wavelength side of the visible solar spectrum. This marks the discovery of UV radiation from the Sun. The
? coronal line. Most of the wavelengths bidden Fe XII 1242 A of the lines are now known with an uncertainty of a few millia? ngstrHm or better (cf., Kelly, 1987), but are given here as rounded values suTcient to identify the lines. ? Only slight limb brightening was found for He II 304 A (Burton and Wilson, 1965), but there is limb brightening of ? relative to H I Ly (Johnson et al., 1958) indiO VI 1032 A cating that the opacity of the solar plasma is not signi
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Table 2 Selection of VUV spectral observations and emission line lists ? Wavelength rangea (A)
? Spectral resolutionb (A)
Solar targetc
Referenced
977–1892 1000 –3000 84 –1216 500 –1850 765 –2600 250 –1300 104 –1892 33–500 977–2802 300 –2803 139 –794 304–1394 200 – 400 1052–2185 51–1270 60 –385 300 – 450 950 –2000 200 –500 50 –300 280 –1370 150 –872 66 –171 1200 –2100 66 –171 164 –765 1175 –1940 1175 –1940 1175 –2100 171– 630 975 –3000 280 –1350 171– 630 977–1936 1175 –1965 303–1343 1175 –1710 160 –1344 1190 –1730 914 –1177 1190 –1730 171– 448 657–1176 465 –1609 308– 633 307– 632 500 –1600 465 –1609
1.0 0.25 – 0.5 0.4 0.6 1.0 1.5 –3.0 — 0.1 0.4 0.3 0.1 2.0 0.2 0.16 0.25 0.04 0.2– 0.3 0.3 0.2 0.26 1.8 0.17 0.11 0.4 – 0.5 — 0.06 0.06 0.06 0.07 0.2 0.05 1.6 0.1 0.06 0.06 1.6 0.05 0.7 0.05 0.06 0.05 ¡0:08 0.044, (0.022) 0.044, (0.022) 0.12 0.3– 0.6 0.044, (0.022) 0.044, (0.022)
QS, limb Disk (QS) Disk (AR) Disk (AR) Disk (AR) Disk QS QS, corona Corona Limb AR QS Flare Eclipse AR AR Flare Limb Flare Disk AR, QS Disk, limb Flare Disk, limb Flare Disk (QS) Limb (QS) Limb (CH) QS Flare Corona Various Flare Flare Corona Sunspot plume Various QS, limb Various Limb, Hare Various AR Limb Corona QS Disk Flare Various
Johnson et al. (1958) Behring et al. (1958) Violett and Rense (1959) Purcell et al. (1960) Detwiler et al. (1961) Hall et al. (1963) Pottasch (1964) Tousey et al. (1965) Burton et al. (1967) Burton and Ridgeley (1970) Freeman and Jones (1970) Dupree and Reeves (1971) Widing et al. (1971) Gabriel et al. (1971) Heroux et al. (1972) Behring et al. (1972) Purcell and Widing (1972) Ridgeley and Burton (1972) Cowan and Widing (1973) Malinovsky and Heroux (1973) Dupree et al. (1973) Firth et al. (1974) Kastner et al. (1974) Samain et al. (1975) Fawcett and Cowan (1975) Behring et al. (1976) Doschek et al. (1976a) Feldman et al. (1976) Kjeldseth-Moe et al. (1976) Sandlin et al. (1976) Sandlin et al. (1977) Vernazza and Reeves (1978) Dere (1978) Cohen et al. (1978) Sandlin and Tousey (1979) Noyes et al. (1985) Sandlin et al. (1986) Lang et al. (1990) Brekke et al. (1991) Feldman and Doschek (1991) Brekke (1993a) Thomas and Neupert (1994) Curdt et al. (1997) Feldman et al. (1997) Brooks et al. (1999) Brekke et al. (2000) Feldman et al. (2000b) Curdt et al. (2001)
a In
some cases, not the full range is covered. all entries fully consistent (FWHM, resolution elements, etc.; second order in parentheses). c AR: active region; QS: quiet Sun; CH: coronal hole. d In chronological order. b Not
are described here.” It goes without saying that the task is not easier now after almost four more decades of research in this
The analyses of the spectra obtained resulted in line lists, in which more and more of the VUV emission lines of neutral and ionized elements could be correctly identi
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Table 2 contains a collection of such lists. The spectroscopic designations of all prominent lines of the quiet-Sun spectrum have been determined by now, but many fainter lines, in particular in spectra of sunspots, Hares and the corona, are
Fig. 3.
still awaiting their identi
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spectra found in the VUV range is made in Fig. 3. It can be seen that the spectra of abundant elements were observed much earlier (larger symbols above the diagonal in the lower panel) than those of less abundant ones. Solar lines of Li, Be, B, Sc, and V are missing altogether in VUV spectra obtained with current instrumentation. Many of such lines, and others are, however, observed outside the range ? 6 6 2000 A ? considered here. It should be pointed 100 A out that the
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observed, and documented in the references cited in Table 2 and the caption of Fig. 3. As one would expect from the diFerent physical conditions prevailing in the quiet Sun, in active regions or in the corona, the sets of lines observed from these sources are not the same, although there is a substantial overlap. FUV observations of stars, not a topic of this review, can provide insight into totally diFerent plasma conditions, albeit without spatial resolution. Solar-type stars show remarkably similar spectra compared with the Sun, a G2 V star (Evans et al., 1975; Ayres et al., 1983; Ayres, 2000; Curdt et al., 2001). The Lyman series of H I and He II, the only VUV hydrogen-like spectra in Fig. 3, and the corresponding continua are of special interest in the study of the temperature structure of the chromosphere. The He II Lyman lines up to ? are contained in a list by Behring et al. Ly at 231:444 A (1976), and an analysis of the He II continuum can be found in Linsky et al. (1976). The H I Lyman lines have been ? for various observed by SUMER up to Ly 20 at 913:87 A solar features together with the He II Balmer lines up to Ba ? (Wilhelm et al., 1997). It must be noted that 22 at 917:74 A no deuterium lines could be seen, which should appear close to the even-numbered He II lines, if the abundance of deuterium were higher. Purcell and Tousey (1960) mentioned ? that neither the He II Balmer lines at 1215.09 and 1215:17 A ? nor the deuterium line at 1215:34 A were detectable in the wing of H I Ly, which is supported by the pro
←−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−− Fig. 3. History of the identi
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? could be measured wavelength of the Na X line at 1111:77 A for the
plasma with electron temperatures of at least several hundred thousand kelvin (Grotrian, 1939; EdlWen, 1943). Although the forbidden lines were observed in the visible wavelength range, their identi
2.2.1. Electron density Gabriel and Jordan (1969b) suggested using the line ratios observed in helium-like transitions for a determination of the electron density in the solar atmosphere. The method depends on the long lifetime of the 2 3 P and 2 3 S terms against radiative decay, and deactivation processes thus can also be aFected by electron collisions. In general, the study of collisional de-excitation of metastable levels plays a rˆole in spectroscopic measurements of the electron density, and can be applied to other iso-electronic sequences as well. This technique is not restricted to lines produced by the same ion, however is much more accurate if line pairs of the same ion can be used, because then there is no need to consider the abundances of ions of diFerent elements or ionization stages. Of particular importance for an understanding of the undisturbed solar atmosphere is the electron density at the base of coronal holes—areas of very low plasma density and source regions of the fast solar wind (Krieger et al., 1973). The forbidden Si VIII lines originating within the ground con
2.2.2.1. Contribution functions Skylab and SOHO observations of the solar chromospheric network in diFerent lines emitted by various ions clearly showed similar structures, which are changing, however, with increasing ionization stages (Reeves, 1976; Lemaire et al., 1997). The reason for this change is the increase in the formation temperature. In ionization equilibrium the electron temperature determines the ionization fractions of the elements in the solar atmosphere (cf., House, 1964; Jordan, 1969). With data from more recent evaluations of Arnaud and RothenHug (1985), Arnaud and Raymond (1992) and Mazzotta et al. (1998), the ionic fractions versus temperature of ions can be obtained as shown in Fig. 4a for some examples. The most signi
2.2.2. Electron temperature The identi
where the
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Fig. 4. (a) Dependence of the ionic fractions of some species on the electron temperature is shown in the upper panel. (b) Relative emissivities versus electron temperature of VUV lines emitted by the selected ions are given in the lower panel. The emissivity in each line is shown as the normalized contribution function. Note the high-temperature extensions of the contribution functions of lines from the sodium-like and lithium-like ions Mg+ , He+ , O5+ , and Ne7+ resulting from dielectronic recombination of neon-like and helium-like ions. Ionization equilibrium was assumed in calculating the ionic fractions (Arnaud and RothenHug, 1985; Arnaud and Raymond, 1992; Mazzotta et al., 1998). The helium and neon curves are shown for two diFerent evaluations to demonstrate the good agreement between them.
the ground con
2.2.3. E?ective ion temperatures As outlined in the introduction, the Doppler formula allows a conversion of measured line shifts or pro
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increase to more than two million kelvin. Waldmeier (1941) later con
observe prominences even without eclipse using spectroscopic means. This method had theoretically been studied by Joseph Norman Lockyer, who independently observed emission lines from prominences in October 1868 (Lockyer, 1868), and later suggested a new element, helium, as the source of the yellow emission line, which could not be matched to any known line on Earth. Only in 1895, helium was found as a terrestrial element by William Ramsay (Lockyer, 1896). After hydrogen, helium is the second most abundant element in the Sun, but the spectroscopic determination of its abundance (relative to hydrogen) is still a major problem, because of the high transition energy of about 20 eV to the
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2.3.1. Transition-region downAows An average red shift of most of the transition-region lines (Si II, C IV, O IV, and N V) was observed in many quiet-Sun disk spectra from the S082-B spectrograph on Skylab by Doschek et al. (1976b). Maximum shifts imply downHows of 15 km s−1 . They are observed in network boundaries. Should these shifts indicate real downHows, then the corona could only support these for a few minutes unless there is some upHow at other temperatures. No shifts in cell interiors were found, although the statistics were poor in this respect. Pneuman and Kopp (1978) presented evidence from the
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parison with solar wind parameters at the Earth. Rottman et al. (1982) observed maximum blue shifts in coronal holes ? relative to quiet-Sun regions for the O V and Mg X 625 A lines, which are consistent with outHow speeds in coronal holes of 7 km s−1 and 12 km s−1 , respectively, if the remainder of the solar disk is at rest. Warren et al. (1997) with SUMER observations also found a blue shift of about 15 km s−1 in coronal holes relative to quiet-Sun regions. Concentrated outHow of Ne7+ ions could be identi
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of the outer convection zone. The network is also visible in the He I and He II lines as well as in all transition-region lines formed between 20 000 and 400 000 K. It fades out at coronal temperatures above 600 000 K. First multi-wavelength observations in the VUV were reported by Reeves et al. (1974) and Reeves (1976), who showed that the contrast between network lanes and cell interiors had a maximum just under 200 000 K. A two-dimensional model of the chromosphere and corona by Gabriel (1976) was able to describe the observations. This model placed the primary transition region on expanding magnetic Hux tubes above the network lanes and a thin secondary transition region above the cells. Huber et al. (1974) had demonstrated that the network is remarkably similar in quiet-Sun areas and coronal holes, but, as seen on the limb, had a greater vertical extension in holes. Similar results were obtained by Feldman et al. (1975). Doschek et al. (1975b) found evidence for either an extended transition region or a need for inhomogeneous models. Without adequate spatial resolution to identify the network, OSO-4 observations had indicated before that the radiances of transition-region lines and continua did not decrease signi
wavelength shifts which seem to vary with the line radiance, a result which could be con
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Much higher LOS velocities of up to 60 km s−1 were observed in cool active region loops by CDS on SOHO (Brekke et al., 1997b). 4.2. Sunspots A Doppler-shifted Mg IX line was observed by SERTS over a large sunspot relative to the surroundings indicating an outHow with ≈14 km s−1 in the low corona (Neupert et al., 1992b). Sunspots are, in general, not prominent features in the VUV spectral range, but sunspot plumes are. The latter are relatively cool structures guided by magnetic
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area, and solid angle) is of primary importance. The latter quantity does not depend on the distance from which the observation is performed (cf., Wilhelm, 2002a). Hinteregger and Hall (1969) and Hall and Hinteregger (1970) used the EUV spectrometer on OSO 3 to study the irradiances of emission lines (not all fully spectrally resolved) ? and their variin the wavelength range from 270 to 1310 A ation with solar rotation. It was found that the lines from highly charged ions, in general with shorter wavelengths, exhibited larger variations than those of species in lower ionization stages. Solar EUV irradiances in the range from ? have been observed by a grating spectrome160 to 1060 A ter on the Aeros A satellite in 1973 together with the impact on the Earth’s atmosphere (Schmidtke et al., 1977). The SOLSTICE instrument on UARS (Rottman et al., 1993) underwent a thorough laboratory calibration traceable to primary radiometric standards. This was followed by a relative comparison of the Sun with stars to achieve a long-term precision for the solar irradiances in orbit (Woods et al., 1993). Data from this instrument were analysed by London et al. (1993). Solar rotation and sunspot cycle effects were studied and their dependence on the wavelength of the radiation was investigated con
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6. Summary and conclusions Many important aspects of the solar upper atmosphere accessible to VUV investigations have not been covered in this review; they include a detailed discussion of elemental abundances, prominences,
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K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189 Behring, W.E., 1970. A spectrometer for observations of the solar extreme ultraviolet from the OSO-I satellite. Applied Optics 9, 1006–1013. Behring, W.E., McAllister, H., Rense, W.A., 1958. Ultraviolet emission lines in the solar spectrum. Astrophysical Journal 127, 676–679. Behring, W.E., Cohen, L., Feldman, U., 1972. The solar spectrum: wavelengths and identi
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www.elsevier.com/locate/jastp
Past and recent observations of the solar upper atmosphere at vacuum-ultraviolet wavelengths Klaus Wilhelm∗ Max-Planck-Institut fur Aeronomie (MPAE), Max-Planck-Strae 2, 37191 Katlenburg-Lindau, Germany Received 4 December 2001; received in revised form 25 June 2002; accepted 27 September 2002
Abstract Our understanding of the high-temperature solar atmosphere is to a large extent based on spectroscopic observations of emission lines and continuum radiation in the vacuum-ultraviolet (VUV) wavelength range of the electromagnetic spectrum. In addition, important contributions stem from soft X-ray measurements. The VUV radiation is produced by transitions of atoms and ions, or to some extent, of molecules. The atomic and ionic emission lines have formation temperatures between 10 000 K and several million kelvin, representative of the chromosphere, the transition region and the corona. The molecular lines and the continua originate in cooler regions of the Sun. Radiation at VUV wavelengths is strongly absorbed by the Earth’s atmosphere leading to important geophysical processes at high altitudes. In our context it means that this radiation can only be detected with instruments on sounding rockets and spacecraft above the atmosphere. Detailed studies of the spectral radiances together with atomic physics data furnish information on the electron density and temperature of the solar atmosphere, as well as on elemental abundances, whereas Doppler line-shift measurements show bulk plasma motions, turbulence, and ion temperatures. Highlights of the research in this
1. Introduction In this review the great signi
∗ Corresponding author. Tel.: +49-5556-979-423; fax: +495556-979-240. E-mail address: [email protected] (K. Wilhelm).
or spacecraft at altitudes above at least 150 km. This absorption, in turn, leads to important processes in the Earth’s ionosphere, including photoionization of N2 , O2 , NO, and O ? reat wavelengths shorter than 796, 1026, 1340, and 911 A, ? spectively, as well as photodissociation of N2 below 1270 A ? Consequently, a detailed knowledge and O2 below 2422 A. of the spectral irradiance of the solar VUV radiation is also essential for a quantitative treatment of the thermosphere. In this context, it is of special importance that the O2 absorp? exactly tion cross section has a deep minimum at 1216 A, at the wavelength of the strong and variable H I Ly line. This line is therefore absorbed at the 85 km level and does not signi
c 2002 Elsevier Science Ltd. All rights reserved. 1364-6826/03/$ - see front matter PII: S 1 3 6 4 - 6 8 2 6 ( 0 2 ) 0 0 2 8 5 - 7
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Fig. 1. Spectral irradiance at 1 ua (astronomical unit) of the solar H I Ly pro
example might help to clarify this point. After the H I Ly line was
work accompanying the observational progress. As far as the observational aspects are concerned, there will be a bias towards the earliest and most recent results. There is just not enough space to present the complete development over more than
K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
the top of the photosphere, in perfect agreement with the value of 4430 K ±50 K found by Samain et al. (1975) from ? Most of the emission lines are the continuum around 1580 A. produced above this temperature minimum and stem from the chromosphere, transition region or corona of the solar atmosphere. The
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approximately v = 1 km s−1 , which are related to wavelengths by the Doppler formula v = c( − o )=o (where c is the speed of light and o is the rest wavelength of a speci
Spacecraftb
Rocket Rockets Rockets Solrad 1 Rocket Rocket OSO 1 Rocket Rocket OSO 3 OSO 3 OSO 4 OSO 4 Rocket OSO 6 OSO 7 Rocket Rocket Skylab/ATM Skylab/ATM Skylab/ATM Rocket OSO 8 OSO 8 Rockets Rockets Rockets SMM
Time
1955 1958/1959 1959 –1961 1960 1960 1961 1962 1963 1965 1967 1967 1967 1967 1969 1969 –1972 1971–1973 1973 1973 1973/1974 1973/1974 1973/1974 1974 1975 –1978 1975 –1978 1975 –1987 1979 –1982 1979/1980 1980 –1989
Telescope, grating (NI) Grating spectrometer (GI) Photoelectric detector (GI) Ly photometer Double dispersion (NI) GI with aluminium
Instrumentc 900 –3000 84 –1216 250 –1300 1050 –1350 500 –1550 170 –700 10 – 400 257–269 180 –2950 260 –1300 20 – 400 300 –1400 304, 1216 60 –385 285 –1385 170 – 400 200 –700 1200 –2100 280 –1340 171– 630 1150 –3940 584 1200 –2000 1025 –1216 1175 –1710 1216 –2200 605 – 633 1150 –3600
Wavelength ? ranged (A) 1 0.4 1.5 –3.0 Bandpass 0.1– 0.2 0.1– 0.2 0.86 — 0.4 2 0.6 1.6 — 0.04 3 0.8 0.02 0.4 – 0.5 1.6 0.027 0.05 15 0.02 0.02– 0.06 0.05 Bandpasses Pixel: (0.028) FWHM: 0.04, (0.02)
Spectral ? resolutione (A) Disk Disk Disk Disk Disk Disk Disk 60 4 Disk Disk 60 × 60 Disk Disk 35 × 35 10 × 20 20 7–10 Slit: 5 × 5 2–3 ... 10 2 × 60 Disk Slit: 2:5 × 3 2 0:5 × 0:8 61 Slit: 18 × 54 62
Angular resolutionf ( )
Jursa et al. (1955) Violett and Rense (1959) Hall et al. (1963) Kreplin et al. (1962) Detwiler et al. (1960) Austin et al. (1962) Behring (1970) Purcell et al. (1964) Burton et al. (1967) Hinteregger and Hall (1969) Neupert et al. (1968) Reeves and Parkinson (1970) Bowles et al. (1968) Behring et al. (1972) Huber et al. (1973) Underwood and Neupert (1974) Cushman and Rense (1976, 1977) Samain et al. (1975) Reeves et al. (1977) Tousey et al. (1977) Bartoe et al. (1977) Maloy et al. (1978) Bruner (1977) Bonnet et al. (1978) Brueckner and Bartoe (1983) Bonnet et al. (1980) Rottman et al. (1982) Woodgate et al. (1980)
Referenceg
Table 1 Selected solar VUV instrumentsa Hown on spacecraft (sub-orbital rockets, satellites, or space probes) and some of their typical performance characteristics
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Spacelab 2 Spacelab 2 Rocket Rockets Rockets Rockets Rockets Yohkoh UARS UARS Spartan 201 SOHO SOHO SOHO SOHO SOHO Rocket TRACE
CHASE (GI) HRTS (NI) Multi-layer telescopes EGS (NI) SERTS (GI) NIXT MSSTA SXT SUSIM (NI) SOLSTICE (NI) UCS (NI) SUMER (NI) CDS (GI,NI) EIT UVCS (NI) SEM (He II) XDT (Fe XIV) Multi-layer telescope
160 –1344 1176 –1700 44 –256 250 –1200 170 – 450 63.5 44 –1548 3– 45 1150 – 4100 1190 – 4200 1032, 1037, 1216 465 –1610 151–785 171–304 492–1287 304, 170 –700 205 –218 170 –1700
0.7 0.05 Bandpasses 1.7 FWHM: 0.05 – 0.08 FWHM: 1.4 Bandpasses Bandpass 10 1–2 Pixel: 0.25 Pixel: 0.044, (0.022) 0.3 Bandpasses 0.09 – 0.14 80, Bandpass ? red/blue 211:3 A: Bandpasses
15 1 1–1.5 Disk 7 0.75 ¡ 0:75 FWHM: 3 Disk Disk Corona: 30 × 150 1–2 4–6 2.5 –5 Corona: 7 Disk 5.2 Pixel: 0.5
Lang et al. (1990) Bartoe and Brueckner (1975) Walker et al. (1988) Woods and Rottman (1990) Neupert et al. (1992a) Golub et al. (1990) Hoover et al. (1991) Tsuneta et al. (1991) Brueckner et al. (1993) Rottman et al. (1993) Kohl et al. (1994) Wilhelm et al. (1995) Harrison et al. (1995) DelaboudiniSere et al. (1995) Kohl et al. (1995) Hovestadt et al. (1995) Sakao et al. (1999) Handy et al. (1999)
b Spacecraft
soft X-ray instruments are included. and spacecraft subsystem abbreviations: ATM, Apollo Telescope Mount; OSO, Orbiting Solar Observatory; SMM, Solar Maximum Mission; SOHO, SOlar and Heliospheric Observatory; TRACE, Transition Region and Coronal Explorer; UARS, Upper Atmosphere Research Satellite. c Instrument abbreviations: GI, Grazing incidence; NI, Normal incidence; CDS, Coronal Diagnostic Spectrometer; CHASE, Coronal Helium Abundance Spacelab Experiment; EGS, EUV Grating Spectrograph; EIT, Extreme-ultraviolet Imaging Telescope; HRTS, High Resolution Telescope and Spectrometer; MCS, MultiChannel Spectrometer; MSSTA, Multi-Spectral Solar Telescope Array; NIXT, Normal Incidence X-ray Telescope; SEM, Solar EUV Monitor; SERTS, Solar EUV Research Telescope and Spectrograph; SOLSTICE, SOLar-STellar Irradiance Comparison Experiment; SUMER, Solar Ultraviolet Measurements of Emitted Radiation; SUSIM, Solar Ultraviolet Spectral Irradiance Monitor; SXT, Soft X-ray Telescope; TRC, Transition Region Camera; UCS, Ultraviolet Coronal Spectrometer; UVCS, UltraViolet Coronagraph Spectrometer; UVS, UltraViolet Spectrometer; UVSP, UltraViolet Spectrometer and Polarimeter; XDT, XUV Doppler Telescope. d In some cases, not the full range is actually covered. e Spectral resolution and resolution elements are not distinguished systematically. Second-order values in parentheses. FWHM: Full-width at half-maximum. f Spatial resolution, resolution elements, or slit dimensions. g References with instrumental details.
a Some
1985 1985 1987 1988–1994 1989 – 1989 – 1991–1994 1991–2002 1991–2001 1991–2001 1993– 1995 – 1995 – 1995 – 1995 – 1995 – 1998 1998– K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189 171
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Fig. 2. Spectral radiance of a quiet-Sun region near the centre of the disk on 12 August 1996. The SUMER grating spectrometer superimposes spectra of the
2. Emission line spectra 2.1. Spectral line identi>cations Johann Wilhelm Ritter observed in 1801 the decomposition of silver chloride on the short-wavelength side of the visible solar spectrum. This marks the discovery of UV radiation from the Sun. The
? coronal line. Most of the wavelengths bidden Fe XII 1242 A of the lines are now known with an uncertainty of a few millia? ngstrHm or better (cf., Kelly, 1987), but are given here as rounded values suTcient to identify the lines. ? Only slight limb brightening was found for He II 304 A (Burton and Wilson, 1965), but there is limb brightening of ? relative to H I Ly (Johnson et al., 1958) indiO VI 1032 A cating that the opacity of the solar plasma is not signi
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Table 2 Selection of VUV spectral observations and emission line lists ? Wavelength rangea (A)
? Spectral resolutionb (A)
Solar targetc
Referenced
977–1892 1000 –3000 84 –1216 500 –1850 765 –2600 250 –1300 104 –1892 33–500 977–2802 300 –2803 139 –794 304–1394 200 – 400 1052–2185 51–1270 60 –385 300 – 450 950 –2000 200 –500 50 –300 280 –1370 150 –872 66 –171 1200 –2100 66 –171 164 –765 1175 –1940 1175 –1940 1175 –2100 171– 630 975 –3000 280 –1350 171– 630 977–1936 1175 –1965 303–1343 1175 –1710 160 –1344 1190 –1730 914 –1177 1190 –1730 171– 448 657–1176 465 –1609 308– 633 307– 632 500 –1600 465 –1609
1.0 0.25 – 0.5 0.4 0.6 1.0 1.5 –3.0 — 0.1 0.4 0.3 0.1 2.0 0.2 0.16 0.25 0.04 0.2– 0.3 0.3 0.2 0.26 1.8 0.17 0.11 0.4 – 0.5 — 0.06 0.06 0.06 0.07 0.2 0.05 1.6 0.1 0.06 0.06 1.6 0.05 0.7 0.05 0.06 0.05 ¡0:08 0.044, (0.022) 0.044, (0.022) 0.12 0.3– 0.6 0.044, (0.022) 0.044, (0.022)
QS, limb Disk (QS) Disk (AR) Disk (AR) Disk (AR) Disk QS QS, corona Corona Limb AR QS Flare Eclipse AR AR Flare Limb Flare Disk AR, QS Disk, limb Flare Disk, limb Flare Disk (QS) Limb (QS) Limb (CH) QS Flare Corona Various Flare Flare Corona Sunspot plume Various QS, limb Various Limb, Hare Various AR Limb Corona QS Disk Flare Various
Johnson et al. (1958) Behring et al. (1958) Violett and Rense (1959) Purcell et al. (1960) Detwiler et al. (1961) Hall et al. (1963) Pottasch (1964) Tousey et al. (1965) Burton et al. (1967) Burton and Ridgeley (1970) Freeman and Jones (1970) Dupree and Reeves (1971) Widing et al. (1971) Gabriel et al. (1971) Heroux et al. (1972) Behring et al. (1972) Purcell and Widing (1972) Ridgeley and Burton (1972) Cowan and Widing (1973) Malinovsky and Heroux (1973) Dupree et al. (1973) Firth et al. (1974) Kastner et al. (1974) Samain et al. (1975) Fawcett and Cowan (1975) Behring et al. (1976) Doschek et al. (1976a) Feldman et al. (1976) Kjeldseth-Moe et al. (1976) Sandlin et al. (1976) Sandlin et al. (1977) Vernazza and Reeves (1978) Dere (1978) Cohen et al. (1978) Sandlin and Tousey (1979) Noyes et al. (1985) Sandlin et al. (1986) Lang et al. (1990) Brekke et al. (1991) Feldman and Doschek (1991) Brekke (1993a) Thomas and Neupert (1994) Curdt et al. (1997) Feldman et al. (1997) Brooks et al. (1999) Brekke et al. (2000) Feldman et al. (2000b) Curdt et al. (2001)
a In
some cases, not the full range is covered. all entries fully consistent (FWHM, resolution elements, etc.; second order in parentheses). c AR: active region; QS: quiet Sun; CH: coronal hole. d In chronological order. b Not
are described here.” It goes without saying that the task is not easier now after almost four more decades of research in this
The analyses of the spectra obtained resulted in line lists, in which more and more of the VUV emission lines of neutral and ionized elements could be correctly identi
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Table 2 contains a collection of such lists. The spectroscopic designations of all prominent lines of the quiet-Sun spectrum have been determined by now, but many fainter lines, in particular in spectra of sunspots, Hares and the corona, are
Fig. 3.
still awaiting their identi
K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
spectra found in the VUV range is made in Fig. 3. It can be seen that the spectra of abundant elements were observed much earlier (larger symbols above the diagonal in the lower panel) than those of less abundant ones. Solar lines of Li, Be, B, Sc, and V are missing altogether in VUV spectra obtained with current instrumentation. Many of such lines, and others are, however, observed outside the range ? 6 6 2000 A ? considered here. It should be pointed 100 A out that the
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observed, and documented in the references cited in Table 2 and the caption of Fig. 3. As one would expect from the diFerent physical conditions prevailing in the quiet Sun, in active regions or in the corona, the sets of lines observed from these sources are not the same, although there is a substantial overlap. FUV observations of stars, not a topic of this review, can provide insight into totally diFerent plasma conditions, albeit without spatial resolution. Solar-type stars show remarkably similar spectra compared with the Sun, a G2 V star (Evans et al., 1975; Ayres et al., 1983; Ayres, 2000; Curdt et al., 2001). The Lyman series of H I and He II, the only VUV hydrogen-like spectra in Fig. 3, and the corresponding continua are of special interest in the study of the temperature structure of the chromosphere. The He II Lyman lines up to ? are contained in a list by Behring et al. Ly at 231:444 A (1976), and an analysis of the He II continuum can be found in Linsky et al. (1976). The H I Lyman lines have been ? for various observed by SUMER up to Ly 20 at 913:87 A solar features together with the He II Balmer lines up to Ba ? (Wilhelm et al., 1997). It must be noted that 22 at 917:74 A no deuterium lines could be seen, which should appear close to the even-numbered He II lines, if the abundance of deuterium were higher. Purcell and Tousey (1960) mentioned ? that neither the He II Balmer lines at 1215.09 and 1215:17 A ? nor the deuterium line at 1215:34 A were detectable in the wing of H I Ly, which is supported by the pro
←−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−−− Fig. 3. History of the identi
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? could be measured wavelength of the Na X line at 1111:77 A for the
plasma with electron temperatures of at least several hundred thousand kelvin (Grotrian, 1939; EdlWen, 1943). Although the forbidden lines were observed in the visible wavelength range, their identi
2.2.1. Electron density Gabriel and Jordan (1969b) suggested using the line ratios observed in helium-like transitions for a determination of the electron density in the solar atmosphere. The method depends on the long lifetime of the 2 3 P and 2 3 S terms against radiative decay, and deactivation processes thus can also be aFected by electron collisions. In general, the study of collisional de-excitation of metastable levels plays a rˆole in spectroscopic measurements of the electron density, and can be applied to other iso-electronic sequences as well. This technique is not restricted to lines produced by the same ion, however is much more accurate if line pairs of the same ion can be used, because then there is no need to consider the abundances of ions of diFerent elements or ionization stages. Of particular importance for an understanding of the undisturbed solar atmosphere is the electron density at the base of coronal holes—areas of very low plasma density and source regions of the fast solar wind (Krieger et al., 1973). The forbidden Si VIII lines originating within the ground con
2.2.2.1. Contribution functions Skylab and SOHO observations of the solar chromospheric network in diFerent lines emitted by various ions clearly showed similar structures, which are changing, however, with increasing ionization stages (Reeves, 1976; Lemaire et al., 1997). The reason for this change is the increase in the formation temperature. In ionization equilibrium the electron temperature determines the ionization fractions of the elements in the solar atmosphere (cf., House, 1964; Jordan, 1969). With data from more recent evaluations of Arnaud and RothenHug (1985), Arnaud and Raymond (1992) and Mazzotta et al. (1998), the ionic fractions versus temperature of ions can be obtained as shown in Fig. 4a for some examples. The most signi
2.2.2. Electron temperature The identi
where the
K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
177
Fig. 4. (a) Dependence of the ionic fractions of some species on the electron temperature is shown in the upper panel. (b) Relative emissivities versus electron temperature of VUV lines emitted by the selected ions are given in the lower panel. The emissivity in each line is shown as the normalized contribution function. Note the high-temperature extensions of the contribution functions of lines from the sodium-like and lithium-like ions Mg+ , He+ , O5+ , and Ne7+ resulting from dielectronic recombination of neon-like and helium-like ions. Ionization equilibrium was assumed in calculating the ionic fractions (Arnaud and RothenHug, 1985; Arnaud and Raymond, 1992; Mazzotta et al., 1998). The helium and neon curves are shown for two diFerent evaluations to demonstrate the good agreement between them.
the ground con
2.2.3. E?ective ion temperatures As outlined in the introduction, the Doppler formula allows a conversion of measured line shifts or pro
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K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
increase to more than two million kelvin. Waldmeier (1941) later con
observe prominences even without eclipse using spectroscopic means. This method had theoretically been studied by Joseph Norman Lockyer, who independently observed emission lines from prominences in October 1868 (Lockyer, 1868), and later suggested a new element, helium, as the source of the yellow emission line, which could not be matched to any known line on Earth. Only in 1895, helium was found as a terrestrial element by William Ramsay (Lockyer, 1896). After hydrogen, helium is the second most abundant element in the Sun, but the spectroscopic determination of its abundance (relative to hydrogen) is still a major problem, because of the high transition energy of about 20 eV to the
K. Wilhelm / Journal of Atmospheric and Solar-Terrestrial Physics 65 (2003) 167 – 189
2.3.1. Transition-region downAows An average red shift of most of the transition-region lines (Si II, C IV, O IV, and N V) was observed in many quiet-Sun disk spectra from the S082-B spectrograph on Skylab by Doschek et al. (1976b). Maximum shifts imply downHows of 15 km s−1 . They are observed in network boundaries. Should these shifts indicate real downHows, then the corona could only support these for a few minutes unless there is some upHow at other temperatures. No shifts in cell interiors were found, although the statistics were poor in this respect. Pneuman and Kopp (1978) presented evidence from the
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parison with solar wind parameters at the Earth. Rottman et al. (1982) observed maximum blue shifts in coronal holes ? relative to quiet-Sun regions for the O V and Mg X 625 A lines, which are consistent with outHow speeds in coronal holes of 7 km s−1 and 12 km s−1 , respectively, if the remainder of the solar disk is at rest. Warren et al. (1997) with SUMER observations also found a blue shift of about 15 km s−1 in coronal holes relative to quiet-Sun regions. Concentrated outHow of Ne7+ ions could be identi
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of the outer convection zone. The network is also visible in the He I and He II lines as well as in all transition-region lines formed between 20 000 and 400 000 K. It fades out at coronal temperatures above 600 000 K. First multi-wavelength observations in the VUV were reported by Reeves et al. (1974) and Reeves (1976), who showed that the contrast between network lanes and cell interiors had a maximum just under 200 000 K. A two-dimensional model of the chromosphere and corona by Gabriel (1976) was able to describe the observations. This model placed the primary transition region on expanding magnetic Hux tubes above the network lanes and a thin secondary transition region above the cells. Huber et al. (1974) had demonstrated that the network is remarkably similar in quiet-Sun areas and coronal holes, but, as seen on the limb, had a greater vertical extension in holes. Similar results were obtained by Feldman et al. (1975). Doschek et al. (1975b) found evidence for either an extended transition region or a need for inhomogeneous models. Without adequate spatial resolution to identify the network, OSO-4 observations had indicated before that the radiances of transition-region lines and continua did not decrease signi
wavelength shifts which seem to vary with the line radiance, a result which could be con
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Much higher LOS velocities of up to 60 km s−1 were observed in cool active region loops by CDS on SOHO (Brekke et al., 1997b). 4.2. Sunspots A Doppler-shifted Mg IX line was observed by SERTS over a large sunspot relative to the surroundings indicating an outHow with ≈14 km s−1 in the low corona (Neupert et al., 1992b). Sunspots are, in general, not prominent features in the VUV spectral range, but sunspot plumes are. The latter are relatively cool structures guided by magnetic
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area, and solid angle) is of primary importance. The latter quantity does not depend on the distance from which the observation is performed (cf., Wilhelm, 2002a). Hinteregger and Hall (1969) and Hall and Hinteregger (1970) used the EUV spectrometer on OSO 3 to study the irradiances of emission lines (not all fully spectrally resolved) ? and their variin the wavelength range from 270 to 1310 A ation with solar rotation. It was found that the lines from highly charged ions, in general with shorter wavelengths, exhibited larger variations than those of species in lower ionization stages. Solar EUV irradiances in the range from ? have been observed by a grating spectrome160 to 1060 A ter on the Aeros A satellite in 1973 together with the impact on the Earth’s atmosphere (Schmidtke et al., 1977). The SOLSTICE instrument on UARS (Rottman et al., 1993) underwent a thorough laboratory calibration traceable to primary radiometric standards. This was followed by a relative comparison of the Sun with stars to achieve a long-term precision for the solar irradiances in orbit (Woods et al., 1993). Data from this instrument were analysed by London et al. (1993). Solar rotation and sunspot cycle effects were studied and their dependence on the wavelength of the radiation was investigated con
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6. Summary and conclusions Many important aspects of the solar upper atmosphere accessible to VUV investigations have not been covered in this review; they include a detailed discussion of elemental abundances, prominences,
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