high resolution uv solar spectroscopy

H I G H R E S O L U T I O N U V S O L A R S P E C T R O S C O P Y *

R. M. B O N N E T

Laboratoire de Physique Stellaire et Plangtah'e, Verri~res le Buisson, France

(Received 24 October, 1977)

Abstract. The advantages of high resolution UV spectroscopy for the investigation of the solar atmosphere are stressed while the limitations in the areas of instrumentation and diagnosis are discussed. The recent achievements (made essentially by Skylab, OSO-8 and rocket instruments) are reviewed and discussed.

It is shown that high resolution UV solar spectroscopy has improved our knowledge of the dynamics of the upper layers of the solar atmosphere. Within the :present instrument capabilities the birth of coronal expansion is shown to take place at the top of the transition region. The existence of downward flows over the bright regions of the network is evidenced from redshifts or transition region and chromospheric optically thin lines: velocities as large as 22 km s -1 have been measured in O vI. Short period waves (95 s) have been detected in lines of Si II at chromospheric levels in addition to the well known 300 s and 180 s photospheric and chromospheric oscillations. There is strong evidence that optically thin chromospheric and transition region lines are broadened by a nonthermal velocity component which is maximum at 1.3 x 10 s K and decreases at higher temperatures. This may indicate the presence of unresolved acoustic or magnetohydrodynanfie waves so oftenly set fourth as the source of chromospheric and coronal heating.

Contradictions between the various results are pointed out and discussed. They might be attributed to the different angular resolution of the instruments, a key parameter for future space observations. It is suggested that the Solar Optical Telescope (SOT) and the Grazing Incidence Solar Telescope (GRIST) which are presently under phase A studies at NASA and ESA be considered as a tandem of instruments to fly on Spacelab in the 1980's. Both their angular and spectral resolution appear sufficient to resolve most of the problems under discussion today.

1. Introduction

It is a general trend of astrophysical observations that any progress made in the knowledge of the physics of galaxies, stars, planets, etc. . , proceed from successive improvements in the angular and spectral resolution of the instruments. Solar UV research, following the same trend, went first through a phase of exploration accomplished with rocket and satellite borne in,;truments of moderate resolution. In the mid 1960's with the advances in space tec, hnology, more refined and bigger instruments were flown above the Earth atmosphere.

Although the first high resolution spectroscopic observations of the Sun trace back to 1959, when profiles of La, resolved over the disk were obtained with an instrument built by NRL (Tousey, 1963), the progress made ever since, have not been able yet to provide unique or unambiguous answers to most of the fundamen- tal questions, high resolution solar UV spectroscopy might investigate.

In this paper, we review the most recent and important results achieved with high resolution UV spectrometers launched in the 70's, together, whenever possible, with their impact on solar physics. For practical reasons, and because most of the

* Review presented at the Vth Conference on UV and X-ray Spectroscopy of Astrophysical and Laboratory Plasmas, London, July 4-7, 1977.

Space Science Reviews 21 (1978) 379-409. All Rights Reserved Copyright (~ 1978 by D. Reidel Publishing Company, DordrecJit, Holland

3 8 0 R . M . BONNET

scientific questions that can be studied through high resolution spectroscopy may find an answer in this spectral range, this review will be limited to wavelengths ranging from 300 ~ to 3000 ~ , a domain that can be mainly investigated through the use of normal incidence diffraction grating optics.

In general, it is assumed that a line profile is 'resolved' when the instrumental Full Width at Half Maximum (FWHM) is = ~ the FWHM of the profile (White, 1977). La is the broadest feature between 300 ]~ and 3000 ~ , with a FWHM of 0.5 ~ , and the line can be considered as being resolved with a spectral resolution of 0.05 ~ at 1200 Zk, i.e. )t/A)t = 24 000, which gives A)t = 0.1 ]k at 3000 ~ . Accord- ingly, we consider as 'high spectral resolution observations' between 300 ~ and 3000~ , those made with an instrumental FWHM smaller than 0.1 ~. 'High resolution' is usually tightly associated with high spatial resolution, White (1977). However, we will not restrict this review to high spatial resolution observations only because these are still scarce and because new insights in the physics of the Sun have very often been obtained with the combination of moderate spatial and high spectral resolution. In principle, by resolving individual line profiles, one gains access to information which can be otherwise totally hindered or only, guessed at, in the best cases. However, any contradiction that may appear between results deduced from observations made with different instruments might well be the consequence of insufficient spatial resolution.

Low spatial resolution observations (worse than 3 arc min) are reported in Section 4, while the subsequent sections deal with observations of spatial resolution better than 2 arc rain.

2. Scientific Questions that Can Be Investigated through High Resolution UV Spectroscopy

Figure 1 is an average model of the chromosphere, transition region and corona which shows the temperature rise above the temperature minimum, an evidence that there exists somewhere in the deeper layers of the Sun, a source of atmospheric heating.

Between 300 ~ and 3000 ~ , light is emitted in the layers represented in the model of Figure 1.

In order to give a broad idea of the importance of each region of the model for UV emission, we schematize in Table I the correspondence between wavelength and temperature. This correspondence is very schematic and we may often find spectral lines in any of these spectral ranges, which originate from either the chromosphere, the transition region or the corona: for example, the line of C Iv 1548 ~ originates in the transition region.

The thermodynamic structure of the atmosphere is in fact strongly dependent upon its state of motion and conversely (Athay, 1976). Hence, the study of line profiles may provide the necessary information to refine the model and to infer the mechanism(s) of emission of lines and continua.

HIGH RESOLUTION UV SOLAR SPECTROSCOPY 381

p-

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B 1-2

Fig. 1. An average model of the upper solar atmosphere showing the temperature rise in the chromosphere, transition region and corona (the figure is reproduced from Athay, 1976).

TABLE I

Schematic correspondence between wavelengths and temperatures in the solar atmosphere

Wavelength range (in ~) Temperature regime in K

3000-1600 1600-1200 1200-600

A <600

8000 (photosphere)--4150 (minimum) 4150-20 000 (chromosphere) 20 000-106 (transition region) T > 106 (corona)

High spectral resolution is necessary to ident:ify spectral lines, or to resolve the various components of a multiplet. It also in principle provides evidence of macroscopic motions and flows of matter in the solar atmosphere through the observation of displacements of lines from their rest position. Similarly it permits to follow the propagation of periodic or aperiodic dynamical perturbations in the atmosphere. As indicated by White (1977), conservation of mass and energy amplifies considerably any small velocity perturbation existing at the base of the chromosphere, when they reach the transition region and the corona. This is why it is believed that UV observations, below 2000 ~, provide a more powerful and highly resolving depth probe of these layers.

By observing in the UV with instruments of high resolving power one may solve some of the most basic problems of solar and plasma physics, such as:

- the emission mechanisms of spectral lines; - the birth of coronal expansion;

382 R.M. BONNET

- the source(s) of chromospheric and coronal heating; are they waves (propagating, or standing), and if so, where in the atmosphere do they dissipate their energies?;

- the energy balance of the various regions of the quiet Sun, sunspots and active regions;

- the dynamical effects associated with solar flares, etc. . . We shall in the following, analyze what the present state of high spectral resolu-

tion observations allows us to deduce in the investigation of these problems. However, it is important to discuss first what are the limitations of the instrumental and diagnostic techniques involved.

3. Instrumental and Diagnosis Limitations of High Resolution UV Solar Spectroscopy

Various techniques have been used to improve spectral resolution in the UV. They have been reviewed in the specialized literature (e.g. see Tousey, 1963) and excellent descriptions of new techniques have recently been given at the Symposium on New Instrumentation for Space Astronomy in Tel-Aviv (e.g. see Bruner and Bonnet, 1977; Bartoe and Brueckner, 1975, etc...).

We show in Figures 2 to 5 as an example the schematics of the 4 instruments which have provided most of the results on which this review is bearing: the Skylab ATM NRL instruments S-082-B, the two pointed instruments on OSO-8 under the responsibility of the University of Colorado and the Centre National de la Recherche Scientifique (CNRS), and the NRL stigmatic High Resolution Telescope and Spectrometer. Table II summarizes their main technical characteristics.

In general, the instruments make use of diffraction gratings. Echelle spec- trographs have been used, mostly at wavelengths longer than 1200 Zk. Non-classical and original techniques have also been used, such as the absorption cells developed by Delaboudini6re et al. (1975, 1976). They measured the profile of the 584 ~ He i line by means of a cell filled with pure helium. A pressure scanning device varies the bandpass of the cell which absorbs at larger distances from line center as the

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TABLE II

Main characteristics of four high resolution spectrometers

Highest Spectral Spectral Spatial temporal

Group of principal range resolution resolution resolution investigator Mission (A) (/~) (arc sec) (s)

NRL Skylab 970-1950 0.06 2" x 60" 10 (Bartoe et al., 1977) 1950-3940 0.12

University of Colorado OSO-8 1150-1950 0.02 1"• to 0.16 1" • 900" (Bruner, 1977a)

CNRS-LPSP OSO-8 1000-4000 0.02 1"• 1" to 0.16 (Artzner et al., 1977a) 6"• 120"

NRL Rocket 1200-1700 0.05 1"• 1" 40 (Bartoe and Brueckner, 1975)

H I G H R E S O L U T I O N U V S O L A R S P E C T R O S C O P Y 385

pressure increases. The instrumental profile is theoretically the Voigt profile of the 584 A line, thermal broadening by helium atoms being predominent. At normal temperatures, the FWHM is: 2 mZk. The device is presently being improved to allow the observation of other lines.

Although extremely sophisticated, all these instruments suffer however from limitations which affect more or less severely their capabilities of diagnosing the solar atmosphere. Since the ultimate goal is to achieve simultaneously high spectral, angular and temporal resolution, the amount of light contained in a single resolution element is in general very low, particularly in the UV and for faint-optically thin lines which are oftenly the most interesting to study. Within the weight and volume capabilities of present rockets and spacecrafts, the size of telescopes is generally limited, and to keep the signal statistics within reasonable limits neces- sitates that either the temporal or the angular resolution, or both, be deteriorated. Limitations may also come from the instruments themselves. For example those utilizing mechanically ruled gratings generally introduce a large quantity of scat- tered light due to the existence of 'ghosts', and suffer from intrinsic limitations in spectral resolution. The introduction of hollographically ruled gratings strongly improves the situation, as shown on Figure 6 from Kohl and Parkinson (1976).

Limitations are not only instrumental however, and the diagnostics also impose constraints which result in increasing the number of assumptions used in interpret- ing the results.

This is the case for two important problems: (i) The determination of the height(s) of formation of a spectral line in the solar

atmosphere. This is more severe in the case of optically thick lines, as shown by Athay (1976),

and large uncertainties affect the derivation of the height, or range of heights, at which a line is formed. In general it is only possible to discriminate grossly between different layers of the atmosphere.

It is also sometimes difficult to determine unambiguously the temperature of maximum emission of a line. This is particularly troublesome for the derivation of nonthermal doppler widths of lines formed in the corona (see Section 8.1 below), since differing ionization theories may lead to differing temperatures.

(ii) The interpretation of doppler shifts and asymmetries. As shown by Athay (1970), one may interpret in two opposite ways the doppler

shifts and asymmetries of the Ca rI H and K lines by assuming either an upward or a downward mass motion. Again, the location of these motions in the atmosphere is also fairly ambiguous. In general, the relation zlA/A = v/c, where v represents the velocity along the line of sight, is not valid for an optically thick line and is also questionable for optically thin lines. Therefore one must exert extreme caution in assigning a doppler velocity to an observed shift. The following discussions bear in most cases on the use of this relation for optically thin lines. Consequently, we warn the reader that the conclusions one may derive are depending upon the validity of this relation.

386 R.M. BONNET

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gratings (crosses and dashed line), (from Kohl and Parkinson, 1976).

1

H I G H RESOLUTION UV SOLAR SPECTROSCOPY 387

Optically thin lines are supposed to be formed in a region of the model of a few 100 kms. This would in principle allow the study of wave propagation since the wavelengths of the shortest periods expected to be measured in the upper chromosphere and transition region are of the order of 1000 km (White, 1977). Nevertheless, the larger the contribution function and the smaller the period of the waves, the less detectable will be the oscillations.

4. Full Sun and L o w Angular R e s o l u t i o n Observat ions (a few arc min)

Such observations are of interest for exploratory spectroscopy, such as the identification of lines, the first observation of line profiles (e.g. McAllister, 1974; Moe et al., 1976a), etc. . .

Several lists of newly identified lines have been published in the past thanks to high spectral resolution. See for example, Malinovsky and Heroux (1973), Behring et al. (1972, 1976), Moe et al. (1976b), Sandlin etal . (1977).

Recently, Kohl et al. (1977) reported on the first positive evidence for the presence of boron in the Sun provided by photoelectric spectra near 2500 ]k with a resolution of 0.028 .~. A spectral synthesis analysis of the wavelength region near the 2 p P ~ lines of B I at 2496 .772~ and 2497 .723~ yields log AB = 2.6+ 0.3 on the scale where log AH = 12, which is in good agreement with predictions of light nuclide formation by galactic cosmic ray spallation.

One of the poorly known species in the Sun has, for a long time, been helium. For example, the width of the 584/~ line has been a subject of controversy between astrophysicists and aeronomists, the lattter, Donahue and Kumer (1971), Kumar et

al. (1973), claiming that the FWHM being on the order of 0.02 ~ and the former, Milkey et al. (1973) estimating it to be =0.14/~, until it was measured by several independent experiments (Cushman et al., 1975; Feldman and Behring, 1974) and found to be 0.14 ~. Later measurements by DelaboudiniSre and Crifo (1976), using the absorption cell technique led to a value of 0.127/~.

Another enigma concerns the mechanisms of the He Ii emissions, which fall into three categories:

- ionization and excitation by electron collisions; - enhanced collisional excitation due to mixing or diffusion; - excitation of the n = 2 and n = 3 levels following radiative recombination of

the He Ii. Of importance in that respect is the Balmer a line of He u at 1640.4/~. This line

is made of seven components and was observed with high spectral resolution by Feldman et al. (1975a) from Skylab, and Kohl (1977) from a rocket borne instrument with higher spectral resolution (0.029 ~). Kohl finds that both collisional excitation and recombination play a role in the line formation. However, neither the collisional nor the recombination mechanism alone predict line widths that are fully consistent with the observations. The recombination mechanism appears to be enhanced in plages.

3 8 8 R. M. BONNET

5. High Resolut ion Observations of Strong Optically Thick Lines

In the spectral range considered here, we find three strong optically thick lines: - the two h and k resonance lines of Mg II;

3p 2P1/2- 3s 2S3/2 (2803 ~ , h line); 3p 2P1/2-3s ZS3/2 (2796 ~ , k line); and

- the La resonance line of hydrogen (1216/~). The h and k lines were observed by various balloon or rocket borne instruments

(see for example Lemaire and Skumanich, 1973; and Kohl and Parkinson, 1976) with resolutions of =0.03 ~ , and by Skylab (Feldman and Doschek, 1977). These lines show asymmetries between the blue and red emission maxima which are identical to those observed in the Ca ~I H and K lines. Simultaneous measurements of the Ca II and Mg II resonance doublets made with the CNRS instrument on OSO-8 show very tight resemblance (Bonnet et al., 1978). Due to the stronger coupling of their source function to the chromospheric Planck function, to the stronger increase in the Planck function at shorter wavelengths, and to the higher abundance of magnesium, the Mg ix emission is much more prominent than the Ca n emission, and asymmetries appear also with a much higher contrast. These lines were studied simultaneously with Lo~ and L/3 (1025 ~) with the CNRS instrument, whose high angular resolution made it possible to distinguish between cell and network intensities and line shapes, as shown on Figure 7. The ratio between the intensities of the blue maxima is 22.3 for La and 1.3 for Mg n k, in this particular example. The asymmetry between the blue and red peaks is found to be much larger in the Mg Ix lines than in La. This difference can be attributed to smaller velocity gradients at the height of formation of Lo~ (Athay, 1970), see also Section 6 below. As shown by Ayres and Linsky (1976) and Vernazza et al. (1973), the observed distribution of intensity across the wings of the lines indicates that partial redistribution prevails in the line forming process.

Figure 8 compares the center-to-limb variation of the Lo~ and L/3 line profiles as observed with the CNRS instrument. One notices easily the strong increase in the central reversal at the limb. The presence of this reversal in the L/3 line is the consequence of the existence of a temperature plateau (Cuny, 1968) at =20 000 K, which can be seen on Figure 1. This important feature of the temperature model could only be deduced from line profile data,

Other optically thick lines of S im and Si n observed by Skylab have been studied by Nicolas et al. (1976) to improve the knowledge of the chromospheric and transition region model.

6. Radial Mot ions Deduced from Doppler Shifts

6.1. OBSERVATIONS OF THE BIRTH OF CORONAL EXPANSION

A first evidence for a radial outflow of plasma in the solar atmosphere was given by Cushman and Rense (1976). They measured from a rocket the doppler shifts of


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iristrument on OSO-8 (Bonnet et al., 1978). The spatial resolution is l"x 10".

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Ly/B LIMB Fig. 8. L a and LB profiles at the center and at the limb of the Sun (/~ = 0.11) obtained with the CNRS ins t rument on OSO-8, (Bonnet et al., 1978). The spatial resolution is 6" • 2'. The two lines in the wings of L/3 are O I 1304.8 ~ and 1305.9/~, appearing in another order of diffraction. Wavelengths increase

to the left. The horizontal scale gives the grating mechan i sm step number .


lines formed at the top of the transition region and in the corona with a resolution

of 0.02 ~. Figure 1 of their paper shows the displacement of the line of Si xI (303 ~), observed above a region assumed to coincide with a coronal hole. The outward radial flow velocity is equal to 16 km s -1 at 1.8 x 106 K, the temperature of Si xI formation. However, according to Doschek and Feldman (1977) 'there was no coronal hole on the Sun the day Cushman and Rense obtained their spectra at their indicated coordinates'. The region of maximum blue shift corresponded anyway to a region of minimum X-ray emission. Their measurement is nevertheless consistent with subsequent determinations described below.

Recently, Sandlin et al. (1977) reporting on the analysis of forbidden lines formed at temperatures ranging from 5 • 104K to 106 K, indicate an observed average blue shift corresponding to 6 km s -1. Brueckner (1977, p. 260) measured doppler shifts of the Fe xII 1349.40 A line and gives an expansion velocity at the base of the corona which varies between 10 km s -t and 15 km s -t, the higher value corresponding to regions of weak photospheric magnetic fields with unipolar character. Above plages the velocity is 0 a: 4 km s -1.

Neupert (1977) mentions that the Fe xv 284 ~ and Fe xvI 335 ~ lines evidence radial flows of 15-20 km s -I relative to the lines of He u which do not evidence any shift and are used as a rest reference (White, 1977).

Finally, Doschek et al. (1976a) indicate that no shifts could be observed from Skylab for the line of O v at 1218.35 ~ formed at T = 2.2 x 105 K.

We may now deduce a consistent figure from observational data, concerning the outward flows related to the beginning of coronal expansion.

(i) The solar wind cannot be observed at temperatures below 2.2 x 105 K. (ii) Its velocity increases from 7 • 1 km s -t at 106 K to 17 + 3 km s -t at 1.8 x

106 K, and may vary within factors of at least two for regions of differing magnetic configuration, the regions of opened field corresponding to the highest velocities.

These values look however a little bit low when compared with those given by theories of the solar wind (Athay, 1976).

6 . 2 . E V I D E N C E OF PERSISTENT FLOWS R E S O L V E D ON T H E DISK

The picture one derives from results analysed to date is far from being homo- geneous and most of the time, confusing, forcing us to distinguish between the various regions observed on the disk.

6.2.1. The Quiet Sun Ne twork and SupergranuIation

A commonly accepted feature of the network is the presence of persistent downflows of material. This was first noticed by Doschek et al. (1976a) from the analysis of the positions of Si IV, C IV and O iv lines and later confirmed by Shine et

al. (1976), Lites et al. (1976), Bruner (1977b), and Brueckner (1977, p. 104). All measurements were made in optically thin lines, relative to lines originating in the low chromosphere or near the temperature minimum, which are used as reference. The observed velocities range from 2 to more than 20 km s -~ depending upon the

392 R . M . B O N N E T

line used: Lemaire et al. (1977) measure 22 km s -1 in O vi 1032 ,~ formed at 300 000 K. There is a tendency for the velocity to be lower for lines formed at lower temperatures. Recently, November et al. (1976) confirmed that the down flows in Si iv and C IV are always observed in regions of the network where the magnetic field is high. This is in accord with the previous observations of Lites et aI.

(1976) and Lemaire et al. (1977) that the stronger redshifts correspond to the brightest regions of the network.

There is more evidence to-day that these downward flows are not an average property of the network and the high angular observations of Brueckner (1977, p. 104) show that in some areas, blueshifts of 10kms -1 are also observed. However, the dominant shifts are to the red.

Over cells, Brueckner finds upward velocities of 4 to 10 km s -1 in La and transition zone lines but not in cooler chromospheric lines, while Doschek et al.

(1976a) found no shifts at all. These latter measurements were made however with a resolution of 2" • 60" and this may strongly weaken the possibilities of observing shifts if they are of small spatial extension.

These results are quite puzzling in general because mass conservation requires that balance be reached between upward and downward motions and because downward flows are observed predominantly over regions where spicules are generally observed (Athay, 1977).

At this stage, more observations with higher angular resolution are needed to conclude, but we may advance the following assumptions based on the already

available observational material: (i) Solar gas moving toward the chromosphere is brighter than gas moving away

from it (Withbroe, 1977). This is in good agreement with the existence of asymmetries in the Ca n, Mg ii and La lines if we accept the model proposed by Athay (!970) showing that the bright K2v peaks correspond to predominant downward motions. As shown below in Section 7, the results of Artzner et al. (1977b) from the CNRS instrument on OSO-8 show clear evidence that blue Mg Ii and Ca u peaks correspond to redshifts of the central cores of the Ca n, Mg n and La lines.

(ii) As suggested by Gabriel (1977), spicules are not moving upward as usually admitted and their apparent motion might be a moving luminosity threshold associated with a growing spicule of chromospheric material. The spicule material would then be supplied by condensation from the corona.

The situation is somewhat more confusing when observations are made at the limb where, in principle, one should be able to observe horizontal flows. Shine et al.

(1976) indeed observed a blue shift of 5 km s -1 on both the east and west limbs in the C ~v lines 1548/~ and 1551 ~, which they interpret as the projection along the line of sight of a downward radial velocity of 35 km s -1. It could also be the result of scattering, in the falling material, of photons emitted in the line at the chromospheric level where the matter is supposed to be at rest. Brueckner (1977, p. 260) also reports on shifts observed at the limb that can be interpreted as the projection of outward moving flows.


The limited amount of data available to date clearly claim for more observations with higher angular resolution before any comprehensive picture of horizontal flow patterms at chromospheric and coronal levels be derived.

6.2.2. Sunspots and Associated Features

Small scale, high velocity downward flows have been observed in the transition zone lines above sunspots umbrae. The most spectacular results are those of Brueckner (White, 1977) who measured velocities as high as 100 km s -1 in the region probably corresponding to the sunspot plume. These flows could not be observed in neither the coronal nor the lower chromospheric lines. Bruner et al.

(1976) also observed a similar effect over a sunspot plume with downward radial velocities of 30 km s -1 in C Iv, which persisted over 2 orbits. These observations are consistent with the model proposed by Foukal (1976) in an attempt to explain the energy balance of plumes over sunspots. This model rejects the assumption of hydrostatic equilibrium and suggests that there is a steady state flow of material into the sunspots from the surrounding hot plage areas. If this were true, we might therefore observe reversed shifts in the two, feet of the loop associated with the spot. White (1977) reports on downward flows of 5 to 20 km s -~ measured over plages. Bruns et al. (1976), describing the results obtained with their Salyut-4 instrument (Bruns et al., 1977) show that two parts of a bright and compact flocculus evidence line of sight velocities in opposite direction, with maxima of +35 km s -1 measured in N v 1242.8 ~ and -106 km s -1 in C III 1176.7 ~. They derive a model of the flow of plasma along the loop. Lites et al. (1976) observed blueshift of 6 to 10 km s -1 along the line of sight over an Ha filament.

Here also, more observations are needed evidently before arriving to a better picture of the dynamics of sunspots and associated features.

6.3. O B S E R V A T I O N S OF FLOWS A S S O C I A T E D WITH T R A N S I E N T EVENTS A N D

S O L A R F L A R E S

These observations are important in the chromosphere, where the excitation of flares is oftenly attributed to either flows of high velocity particles or to mass condensation from the corona. Such plasma flows should be detected in the UV lines formed at chromospheric levels and in the transition region.

Although the observational material from both Skylab and OSO-8 which has been published to date is scarce, there is evidence for strong doppler shifts observed during the flare events. The June 15, 1973, flare observed from Skylab is evidenced as a blue shift during the eruptive phase, in lines formed in the transition regions (Doschek et aL, 1977). A rest component and a blue shifted component can be identified in the profiles. The blue shift corresponds to velocities varying between 50 to 84 km s -I depending upon the line. However, no shift of forbidden lines of highly ionized iron, greater than 20 km s -x, could be observed throughout this event (Doschek et al., 1975). The interpretations are ambiguous but do not rule out

394 R . M . B O N N E T

the possibility that continuous plasma ejection occurs over a time interval compar- able to that over which the blue shifted component was observed.

Brueckner et al. (1976) observed a flare event from Skylab on 2 September 1973 and noticed simultaneous blue and red shifts occuring in lines of different ionization potential. Transition zone lines intensity and profile instabilities seem to precede the occurrence of the flare. They show that flare activity is connected with the formation of small loops.

Another event was recently observed with the CNRS instrument on OSO-8 in La, O I 1304 ~, Mg II h and k with a resolution of 40 s in time, and 6" • 2' in space

1200 ~ , , , �9 , , , ,

t k n o : 2 , 1 m

8 0 0

0

FLARE M g e k PIIOFIL|S

�9 �9 " 6. 1 . 1 A , - ; " " "

4 ~

I I I I I I I I ~

Z 11.6

v i , i

16.7

' ' ' ~ - 1.~ A -.~ ' ' "

20.5

, , , u , v �9

2 2 . 2

Fig. 9. T ime evolution of the profile of the M g l I k line during the flare event of 19 April, 1977, as observed with the CNRS ins t rument on OSO-8. The dashed profile refers to the 10" active region blaekground, observed before the occurence of the flare. The extremely large width of the line does not allow to separate visually the various components of the flare. The profiles can be well represented with

two components : one shifted 50 km s -1 to the blue and the other, 20 km s -1 to the red.


(Jouchoux et al., 1977). The OI line formed at the base of the chromosphere showed only weak effects contrary to the La and the Mg II lines which showed intensity enhancements of 15 and 6 respectively. The first evidence of the flare in Mg It profiles appears as a component shifted 20 km s -1 to the red, followed 40 s later, by a weaker component shifted 20 km s -1 to the blue. Figures 9 and 10 from Jouchoux et al. show the evolution of the Mg II k and La profiles as a function of time. They evidence the strong enhancement in both the intensity and width of the lines as well as the presence of several components in the profiles. The strong broadening of the wings even far from line center is also verystriking.

Brunet and Lites (Brunet, 1977b) have found a number of flare like events in the C Iv velocity studies. These events are characterized by a sudden increase in intensity of 5 to 10 fortransition region lines, accompanied by a redshift equivalent

180 �9 . , .

120 1 llmeJ 2 Im

8O

0

........ \

' 13.3 "16~.7

~8.8 :-.42A: ~o.5 ' ~.2-

Fig. 10. Time evolution of the La line profile during the flare event of 19 April, 1977, as observed with the CNRS instrument on OSO-8, See Figure 9 for more details. Notice the strong intensity and width

enhancement of the line together with the appearance of several detailed features in the profile.

3 9 6 R . M . BONNET

to motions greater than 20 km s -1. The increase in intensity can be as large as a factor of 3. No flashes were found showing blue shifts. Lites and Hansen (1977) also report on UV brightening flare like events which they interpret as the chromospheric response to beams of particles or to downward propagating shocks. The data available now are insufficient to rule out definitely any of the models of flares.

All the observations reported here clearly show the need for higher spectral, angular and temporal resolution observations. Spectral resolution is necessary to separate the various components of the flare and angular resolution of at least one arcsec, is necessary to resolve the kernel(s) of enhanced emission. The Solar Maximum Mission to be launched by NASA at the end of this decade may provide this desired opportunity.

7. Observations of Resolved Velocity Fields in the Chromosphere and Transistion Region

The investigation of the heating mechanism(s) of the chromosphere and corona goes through a systematic search for waves, acoustic or magnetohydrodynamic. The

.observed fluctuation of the position of optically thin lines, provided their contribution function is small compared to the 'wavelength' of the searched wave, is usually attributed to the existence of waves. Since, the wavelength is the product of the period by the velocity of propagation, acoustic waves may be easier to observe higher in the atmosphere consequently to the increase in the velocity of sound. However, the waves of ---30 s period which theoretically dissipate their energy in the chromosphere and heat the corresponding layers will be more difficult to observe.

The first systematic attempt to detect chromospheric and transition region velocity oscillations was undertaken by the OSO-8 pointed instruments: they had the required temporal resolution and observed for sufficiently long sequences of time. The University of Colorado instrument had the capability of studying several optically thin lines formed at various heights in the upper atmosphere. Athay and White (1977) report on power spectrum analysis of time series of positions and intensities of the Siu 1815.93 A, 1817.45 ~ lines observed over cells, network, plages and sunspots, with an angular resolution of 1" • 20". They do find broadely spread power peaking at -=-300 s in plages and at ---200 s in supergranule cells. A detailed analysis of individual power spectra shows that oscillations with frequen- cies between 10 and 33 mHz (periods between 33 and 100 s) are relatively common at the height of formation of the Si Ii lines. They also show the first evidence of strong power at short periods (=95 s) in a plage. The velocity amplitudes always rest in the range of 1 km s -1.

Bruner (1977b) mentions that both the 300 s and 95 s oscillations are observed at a fraction of ---50% of the areas examined.

Chipman (1977) also notices very strong power at 300 s in both the velocity and intensity fluctuations of the C u 1336 ~ line formed at around 20 000 K, but finds


little evidence for power in the range 150-200 s. Figure 11 shows his results for a plage region. The peak to peak velocity amplitude is 2 km s -~. In regions of normal or enhanced chromospheric network the 300 s peak in the velocity power spectrum is much broader. Comparison of line and background intensity variations shows upward wave propagation with time delays of 27 s and 70 s for different cases between the 1800 A continuum and the C n 1336 ,~ line intensity variation. These results are consistent with acoustic propagating waves. The same analysis for the C IV 1548 ,~, 1550 .~ lines does not reveal any power at a particular frequency, but only white noise, However, power spectrum analysis of the intensity of lines formed at chromospheric levels, made by the Harvard group on Skylab data

EXPERIMENT FITTED INTENSITY, AVERAGE INTENSITY : 1304

Z 0 TIME (sec)

O 500 ~OOO 15OO 2000

o-I . Ahq l II IIA.J

5 z

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|,, , .

O ~ FREQUENCY

j, z ~

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- - ,)

> ~ -z- BLUE

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o E o . ~ t -

O- (sec)

(MilliHertz)

5266 VELOCITY

TIME (sec] 500 rOOO 1.5OO 2000

V " I " I ,..i

500 3 0 0 200 I 0 0

FREQUENCY (MilliHeri'z)

5 E:~ i-

>'~- 0- :.EE ~ -~-

w r ~e2 2~ ,.o-

Q5-

WIDTH, AVERAGE WIDTH = 22 (KM/sec) TIME ( s e c )

O 500 ioO0 ISO0 2000

o I00

0 I0 ( M i l l i H e r l z )

300 200 r - - - - - 4

FREQUENCY

BACKGROUND INTENSITY, AVERAGE INTENSITY = 227 z TIME ( s e c ) o ,~ 0 500 IO00 1500 2000

} , I h I ,, I ~- o 204

~_~ o

u oc ~ 4 0 0 -

8 2 0 0 -

:~ . . . . i i I(~10

i 10

O- PERIOD (see}

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FREQUENCY (MilliHerlz)

Fig. 11. Power spectrum analysis of the intensity, velocity and width fluctuation of the Cn 1336 resonance line and of the background intensity as observed on 2 September, 1976 at 2316 UT by the

University of Colorado instrument on OSO-8.

3 9 8 R. M~ BONNET

do not evidence power at 300 s, in contradiction with these results (Withbroe, 1977).

The CNRS instrument on OSO-8 also observed oscillations simultaneously in Ca II, Mg II and La but was unable to reveal strong evidence of 300 s oscillations at least in the data samples analyzed to date, Artzner et al. (1977b). However, no power spectrum analysis has yet been made and the results reported here only deal with observations of wave trains, of sometimes more than five periods. Figure 12 shows simultaneous measurements of the position of the Ca II K and Mg t I k

I

i

!

Ca]s K - 200 =

MgTr k

200 s

l

i .oo Bb.oo ='oo.oo ~'zo.oo t',o.oo ,'6o.00 ~o.oo zbo.ool~o.oo - ~ - - ~ o . o o ~'8o.oo 3'o0.00 ~'~o.oo 3',0.00

Fig. 12. Osci l lat ions o f the ccntroid o f the Ca II K and M g II k emission components. Hor izonta l units are mult iples of 10 s (Bonnet eta[., 1978).


centroid, made over the network with a slit of 1" x 10". Wave trains of periods of the

order of 180 s are easy to see. If translated in terms of velocity (with the restrictions discussed in Section 3), the peak to peak amplitude varies between 1.5 and 2.7 km s -1 on this particular example. For the first time, Artzner shows evidence that the displacements associated with these oscillations are also observed at La. He overcame the problem of low signal level by averaging all Lee profiles corresponding to h2v maxima on the one hand and to h2v minima on the other hand. The results are shown on Figure 13. There is evidence that profiles with k2v maxima are reshifted in all 3 lines. The same procedure applied to averaged profiles selected randomly in time does not evidence any such shifts (Figure 14). The estimated amplitude of the shift at La is =3 km s -1. Brueckner (White, 1977), also noticed rapid time changes in the strengths and position of chromospheric lines over intervals as short as 40 s. He also observed different doppler shifts for the various components of a C I multiplet which have very sharp contribution functions and might evidence propagating effects.

The interpretation of these observations is somewhat ambitious at this stage, given the limited set of data yet analyzed. More is expected to come out of the OSO-8 instruments in the future, as data become available. The strong power shown in all Colorado data at 300 s is somewhat puzzling. Recent observations of the fluctuation of the Fe xrv 5303 A line by Tsubaki (1977) show that power is also present at 300 s in the corona. On the other hand Avery (1976) analyzing the fluctuations of the continuum microwave flux at 2.8 cm finds low amplitude quasiperiodic fluctuations at periods of 234s and 150s with a resolution of 2.6 arc min, but no evidence for strong oscillations at periods greater than 250 s.

It is our belief that oscillations in the chromosphere and the corona should be studied with the same resolution at all wavelengths if valuable comparisons and conclusions are to be derived. In the other case, one may find results which look contradictory. It is not impossible indeed that different heating mechanisms play for differing regions of differing magnetic field intensity and geometry, and one should be completely misled in attempting to derive a unique interpretation from sets of data which are intrinsically referring to different regions of strongly different physical conditions.

8. The Analysis of Line Widths and the Search for Unresolved Phenomena

8 . 1 . E V I D E N C E FO R N O N T H E R M A L V E L O C I T I E S

Observations of optically thin line shapes may reflect more easily the distribution of temperature and of unresolved motions, as opposed to optically thick lines whose profiles are function of non linear radiation transfer effects.

All of the profiles formed in the high chromosphere, transition region and low corona have widths exceeding the thermal width corresponding to the ionization and excitation temperature. Many authors, like Boland et al. (1973) have suggested

! J

OS

08

LPS

P

,1

LYM

AN

R

LPH

R

M@K


that the excess width demonstrates the existence of a nonthermal kinetic energy

component which may be due to short period acoustic waves which, as seen previously, cannot be observed, individually, particularly in the case of lines of broad contribution function. Since these waves are the most likely candidates for heating the upper chromosphere, needless to stress the importance of such observations.

Following Feldman et al. (1975b), the FWHM of an optically thin line can be expressed as:

AA2 2 2 2kT i 2 1/2 FWHM= +4(Ln ) ( ,~ /c ) \ - -M-- -+,o] j ,

where Ah 2 is the instrumental FWHM known from laboratory calibration, .~, the wavelength, ~:o, the most probable nonthermal velocity assuming the distribution to be gaussian.

If Ti is assumed to be the electron temperature of maximum ion concentration, one obtains ~:o from the measurement of the line FWHM.

The values for ~:o depend obviously upon Ti which itself depends upon the theories, assumptions and atomic parameters, used in the line formation compu- tations.

With Skylab, OSO-8 and rocket experiments, we do have now a rather good picture of the variation of s with temperature and for various regions of the disk. The most complete set of observations is that obtained by Skylab. On Figure 15 from Moe and Nicolas (1977), s is plotted as a function of the mean temperature of formation. Similar results are reported by Doschek et al. (1976b), Feldman et al.

(1976a) for both limb and disk spectra. They generally agree with those of Boland et al. (1975): ~o increases with increasing Ti up to a maximum of 25 to 30 km s -1 for Ti ~- 10 • 3.105 K and then decreases in the corona. Nicolas et al. (1976) notice a systematic increase of the FWHM of the Si ui 1298.96 ]~ and 1892.03 ~ lines as one proceeds to greater distances above the limb. This may reflect the increase in nonthermal velocities along the line of sight with altitude above the solar surface. The limited angular resolution of the Skylab instrument may also lead to another interpretation, i.e. the superposition of line profiles radiated by each individual element of the atmospheric fine structure in the field of view.

--Fig. 13. The observation of chromospheric oscillations, as observed with the CNRS instrument on OSO-8 (Artzner et aL, 1977a). The upper left quadrant shows the oscillations of the asymmetry of the Mg u k line. Maxima correspond to profiles with strong blue peaks (k2v maxima), and minima to profiles of reversed asymmetry (k2v minima). The intensities for each data point of the Ca u K, Mg rI k and La profiles coinciding to k2v maxima have been added together and an average profile obtained for each line by dividing the added intensities by the total number of profiles, thus giving an average k2v maximum profile (profile A). The same procedure is repeated for profiles coinciding to k2v minima, giving an average k2v minimum profile (profile B). Profiles A and B are plotted in the three other quadrants of the Figure for Ca II K, Mg It k and La. Shaded areas correspond to regions where intensities of profiles B are larger than those of profiles A, evidencing that profile A (kzv maxima) is

redshifted relative to profile B for all 3 lines.

]50

8

LPS

P

f l

~il

t.t

i,l~,l~

t

MC

-K

j ~

~e

,8~

7~

2.~

t ~

6,7

5

~

7~

8,5

~

~

8L

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~

80

4,3

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8

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~,rC CI o r

il 1,0I

4.0 415 5'.0 5'.[5 6.0 6.5 LOG T e

Fig. 15. Nonthermal velocities as 'a function of temperature as observed from Skylab by Moe and Nicolas (1977). Dots and triangles refer to allowed and forbidden transitions respectively.

In view of the interpretation of Boland et al., that nonthermal broadening may be caused by acoustic or MHD waves heating the corona, it is of interest to see whether ~0 has an isotropic distribution. In principle, one assumes isotropy if center and limb values of ~:0 are equal within the experimental errors. Disk and limb comparison were done by Feldman et aL (1976b) and Doschek et al. (1976a) who conclude positively on isotropy for the more intense lines, observed also at Sun center by Boland et al. (1975). There is however a difference between limb and disk profiles for lines of low temperature such as Si IIL Moe and Nicolas (1977) found evidence for a broader component in the wings of strong lines, observed at the limb, contributing 5 to 10% of the intensity, and whose velocities vary between 45 and 75 km s -1.

It is of interest also to measure ~:0 over various regions of the disk. Such detailed. comparisons have been made by both Skylab and OSO-8.

8.1.1. Comparison between Cell and Network

Contradictory results have been obtained by Skylab and OSO-8. Feldman et al. (1976b) claim that they see no differences between cells and network boundaries, while there is increasing evidence from OSO-8 that such differences do exist: White (1977) indicates an increase of 10 to 20% of the FWHM of Si rv and C Iv from cell to network. This has been recently confirmed by Bruner (1977b).

8.1.2. Comparison between Quiet Sun, Coronal Holes and Active Regions

Figure 16, from Doschek and Feldman (1977) shows that there does not seem to be any difference in the widths of coronal lines observed in the quiet Sun and in a hole.

, Fig. 14. Same as Figure 13 when profiles A and B are built out of individual profiles randomly selected in time. The redsbift is nearly absent.

404 R. M. BONNET

I '?

~'E o

200

IOO

8Q

60

40

~0

tO

CH +8"

I (uppg limat)

2

I 0 0 - -

8O

6 O

4 0

CH +8"

2C

~ o 0.25A

CH + 20"

I (upper llmit)

l

Fe XI (1467.1/~)

Q$ + 20"

o 023A

l 0 )ITA

A R - I + I0"

CH +~0"

si wT 0445.8L)

0$ + 20"

O. 26A

o O. 28A

AR-I

O.21A

9~

AR-2 +IO"

AR-~ I + I0"

o r O.ZOA

Fig. 16. Forbidden lines of Si v n i and Fe xI observed outside the limb over the quiet Sun (QS), a coronal hole (CH), and active regions (AR), (from Doschek and Feldman, 1977).

This confirms earlier results of Feldman et al. (1975b, 1976a, 1976b), despite a slight tendency for FWHM to be larger over coronal holes, and implies, according to Withbroe (1977), that the wave flux passing through the transition region is the same for both types of regions.

On the contrary, substantial differences are noticed between active regions and sunspots. Doschek and Feldman (1977) show that the FWHM of coronal lines is 22% smaller in active regions than in the quiet Sun, indicating that nonthermal


motions can be less in regions of strong magnetic fields for lines formed between 104K and 1.7• 106K. On OSO-8, Bonnet et al. (1978) found generally no difference between quiet Sun and active regions in the FWHM of O v1 1032. Over a sunspot, Cheng et al. (1976) have measured FWHM's of lines formed at temperatures ranging from 10 4 tO 2.2 x 105 K. Typical mass motions in the spot are 13 km s -1 at T = 105K compared to 25 km s -a or more in the quiet Sun. They interpret their result as evidence that, if Parker's theory of sunspots is valid, the mode of the hydromagnetic waves is predominantly Alfv6nic. Brueckner (1977, p. 476), with higher angular resolution, shows that reduced line widths are indeed observed over the penumbra of sunspots while no difference is seen above the umbra. Skumanich et aL (1977)with the CNRS instrument observe the FWMH of O vI to be 4 km s -1 greater in the plumes just above sunspots. Obviously this apparent contradiction needs confirmation from more sets of data, corresponding to the same region.

Several measurements of the FWMH of UV lines during the various phases of flares (Brueckner et al., 1976; Doschek et al., 1977; Feldman et al., 1977; Jouchoux et aL, 1977; Bruns et al., 1977) show increased nonthermal broadening, and anisotropies. There is no way of deducing whether these are the result of the mixing of many small scale plasma elements randomly emitting at the time of the flare, or any other cause.

8 .2 . E V I D E N C E FOR LINE ASYMMETRIES

It has been shown by McWhirter and Wilson (1974) and Elzner (1975) that asymmetries and blueshifts should show up in the profiles of transition regions lines as the result of shock propagation and line formation processes in the shock front. The asymmetries should be predominantly to the blue.

Observational evidence for asymmetries exist from several sets of ,observation, see for example Feldman and Behring (1974), Doschek and Feldman (1977), Doschek et al. (1976a, b), Bonnet et al. (1978). Asymmetries are either to the red or to the blue. Boland et al. (1975), Shine et al. (1976) and Brunet (1977b) also observe symmetrical profiles, although the far wings of the lines have a greater intensity than the best gaussian fit for the line center. May be are they seeing the same effect as Moe and Nicolas (1977) who concluded on the existence of a high velocity component in the nonthermally broadened line profiles (see Section 8.1 above).

The situation is therefore not as clear as one would like it to be. To conclude on this section, we may say that the quality and the amount of data,

at the present stage of analysis cannot rule out the possibility that the residual nonthermal velocity and line asymmetries are the manifestation of either acoustic or MHD propagating waves. It is clear that measurements with higher angular and spectral resolution are needed, together probably with measurements of the magnetic field at chromospheric levels with as high an accuracy and a resolution as possible.

406 R . M . BONNET

9. Summary and Conclusions

The overall balance of the past years work in high resolution UV solar spectroscopy is largely positive. Observations have disclosed some of the expected and also new and unexpected results about the dynamics of the chromosphere, transition region and corona.

We do have now a better picture of the solar wind flow and have evidence of a downward flow in the chromospheric network at regions where the magnetic field is strong. The question of horizontal flows is less clear. Sunspots are indeed the field of dynamical phenomena and the assumption of hydrostatic equilibrium appears less and less valid as more observations with increased angular resolution become available.

Oscillatory velocity fields do exist in the chromosphere up to at least a temperature of 20 000 K, but the question of what type of waves and what periods do exist at a given altitude is still to be clarified. For the first time, periods of 95 s have been detected at chromospheric levels. This is the first evidence for the existence of short period waves at these levels. The strong power noticed at 300 s by the University of Colorado Investigators on OSO-8 is puzzling since little or no power at this period is observed in the data of the Harvard ATM group and few wavetrains of periods larger than 200 s have yet been observed with the CNRS instrument on OSO-8.

Evidence of the existence of nonthermal velocity components in the widths of optically thin lines is overwhelming, and their variation with temperature shows a maximum at 1.3• with a tendancy to decrease at coronal levels. Unfortunately, observations have not yet been done with sufficient spectral and angular resolution to discriminate between MHD or acoustic waves as the possible cause of this nontlaermal broadening.

The overall impression on the results obtained so far is that, they are plentiful and of excellent quality in general, but they are still fragmentary in view of diagnosing unambiguously the solar upper atmosphere. The amount of information contained in the available data is really enormous, however their interpretation is not straight-forward and sometimes selfcontradictory, particularly when inter- comparisons are tempted between sets of data obtained with different instruments and different spatial resolution. We do think it important to insist on this last instrumental parameter: the high resolution observations made at NRL disclose new and unexpected aspects of the solar UV spectrum, at the same time they raise new questions. We strongly stress also the importance of having simultaneous observations in the range extending from the minimum to the maximum of temperature, increasing the depth probe of UV observations, releasing the ambi- guity of the location on the disk of observations made at different wavelengths.

Unfortunately, as sad as it may be, no instrument is yet existing which is able to study the Sun at the same time with high temporal, high angular and high spectral resolution. Such an instrument requires the availability of large satellites or large platforms with heavy weight and large volume capabilities. The space Shuttle with


Spacelab will provide such an opportunity in the years to come. In particular, both NASA and ESA (European Space Agency) are presently studying high resolution instruments designed to tackle some of the most important problems outlined above. They are:

(i) the NASA Solar Optical Telescope (SOT), a 1.25 m Gregorian Telescope to study the spectral region extending from 1200 A to the visible with an angular resolution of a fraction of an arcsec, by means of high resolution (A/AA = 10s-106) spectrometers of various types.

(ii) the ESA Grazing Incidence Solar Telescope (GRIST), to study the spectral range 100-1700 .~ with an angular resolution better than 1 arcsec and a spectral resolution of -~ 10 4 to 10 5.

The combination of these two instruments, provided they are co-aligned so that they study exactly the same region of the disk, should provide the first really powerful tool of investigation of the dynamics of the solar atmosphere. This is why we insist on the advantage for solar physicists to have the possibility to fly on and get access to this unique facility in the next decade.

Acknowledgements

This review bears on some new and unpublished results as well as on the information available in the literature. Acknowledgements for making their informations available to the reviewer go to Drs G. Artzner, A. Jouchoux, and P. Lemaire, of the Laboratoire de Physique Stellaire et Plan6taire, to Drs G. Athay, A. Skumanich, and O. R. White of the High Altitude Observatory, to Dr E. C. Bruner, of Lockheed Palo Alto Research Laboratory, Dr E. Chipman, from the University of Colorado, to Drs Doschek and Brueckner of the Naval Research Laboratory, Drs J. Kohl and E. M. Reeves of the Harvard Center for Astrophysics, and to Dr B. Boland of the Rutherford Laboratory (Science Research Council).

The author is also particularly grateful to Drs E. C. Bruner, O. R. White, and G. Brueckner for their useful comments and criticisms.

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408 R. M. BONNET

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