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Infrared coronal lines
Vistas in Astronomy,
Vol. 19, pp. 341-353.PergamonPress, 1976.Printedin Great Britain
INFRARED CORONAL LINES (Observations of Infrared I-Iron XIII] at the 1973 Total Solar Eclipse) JAY M . PASACHOFF a n d DANIEL F. MUZYKA Williams College--Hopkins Observatory, Williamstown, Massachusetts.
IT HAS only been recently that vidicon systems have gained sufficient sensitivity as to allow their use in ground-based infrared studies of the solar corona. This paper reports on the use of a silicon vidicon spectrometer for ground-based eclipse studies of infrared coronal lines. Two forbidden spectral emission lines in the ground state of Fe xm in the near infrared at 10,747 A (3P 1 ---. 3po) and 10,798 ]k (aPE --~ aP1) are the strongest coronal lines in line-to-continuum ratio accessible to ground-based observers. It has been shown that the ratio of the intensity of these lines is insensitive to temperature and provides a good measure of the electron density in corona since one level is radiatively and the other collisionally populated. Simultaneous observations of the 10,747 A and 10,798 A lines therefore provide a sensitive probe of coronal electron density. Both coronagraphic and eclipse observations of these lines have been made but all observations have been severely handicapped by the fact that the lines lie beyond the effective limit of most photographic film and on the upper end of the wavelength limit of the S-1 photocathode of the most commonly employed image tube. At the total solar eclipse of June 30, 1973 in Kenya, we observed the spectral region that included the two lines with a Tektronix J20/TJ20 rapid scanning silicon vidicon spectrometer. We scanned the corona in height from 1.1 to 2.0 RQ. We report here on the observations and their reduction, as well as describe the instrumentation and the theory of the [Fe xIII] lines. INTRODUCTION
Using the coronograph at the Pic du Midi observatory in 1936, Bernard Lyot made a survey of the spectrum of the corona from 3550 to 12,000 A. Conspicuous among the emission lines he found in the coronal spectra were two at 10,747 and 10,798 A that have the two highest line-to-continuum ratios of any of the lines discovered. They showed up "after an extensive trace of the spectrum without coronal radiation" (Lyot, 1939). Edl~n (1942) identified the lines 3 years later as the aP x --~ 3P o (10,747 A) and aP 2 ~ 3P 1 (10,798 A) forbidden transitions of the 3sZ3P2, ground state of Fe xIn, twelvetimes-ionized iron. The importance of these lines was discussed by Firor and Zirin (1962), who showed that the 10,798 A line could not possibly be excited by photoexcitation. This realization and further development of the excitation theory has led to the understanding that one level is basically radiatively and the other collisionally populated, and that a ratio of the line intensities is an accurate tool for measuring electron density in the corona. The theory has been elaborated by Malville (1967), Chevalier and Lambert (1969), Ratier and Rozelot (1972), and Finn and Landman (1973). Though minor corrections in various parameters have been suggested, the theory as set forth by Chevalier and Lambert is now widely accepted. It has been shown both theoretically and observationally that the intensity ratio I(I0,747 ]~)/I(10,798 ]~) is on the order of 7 or less for electron densities much less than 109 c m - 3 and saturates at a value of 2 for densities above that amount. The modern era of observations of these lines began in 1936, when Soviet observers at Pulkovo and Kislovodsk observed the lines with an image tube system (as reported 341
342
JAY M. PASACHOFFand DANIEL F. MUZYKA
by Shkiovsky, 1965). Despite the discovery and identification of the lines in the 1930's, it was 20 years before anyone again observed them. The major problem with observations in those 20 years was that observations could not have been made with any more resolution or structure showing than those made by Lyot--the film limitation on near infrared photography continued. Eclipse observations such as those we and others have since made would have been even more difficult to carry out had the only method of observing been with photographic film. Non-eclipse, coronograph, observations resumed at Climax when observations of 10,747 and 10,798 were made in 1959, subsequently studied by Firor and Zirin (1962), and in 1962 on dates that closely coincided with the New Guinea total eclipse, subsequently studied by Malville (1967) and Zirin (1970). Other non-eclipse observations were reported by W16rick et al. (1963), Dumont and Perche (1964), Fisher and Pope (1971) and Ratier and Rozelot (1972). During the total solar eclipse of 12 November 1966, Byard and Kissell (1971) observed the [Fe xm] lines from the NASA 990 off the southern coast of Brazil. Overall, the observations yield electron densities in the range of l0 s to 109 cm -3, depending upon the height above the limb at which the observations were taken, and on the values of excitation parameters used in the calculations. Other observations of 10,747 A were made to measure line polarization at the 1965 eclipse (Eddy and Malville, 1967) and at the 1966 eclipse (Eddy et al., 1967, 1973). A major problem, as pointed out earlier, is that all observations have been hampered by the low sensitivity of the observing instruments and film. Lyot used Eastman Z-emulsion hypersensitized with ammonia in his spectral line search in the corona. Kurt observed the lines with an electron-optical image converter with an oxygen-cesium cathode. This system had a reported sensitivity of 100 times over the system used by Lyot. A single-stage converter was also used for the Climax observations. The tube, an RCA 6032, had an S-1 photocathode, which has a quantum efficiency at 10,800 A of approximately 0.4%. The converter's output was focused on a phosphor screen to be photographed or viewed. This unit was perhaps slightly more efficient than the unit used by Kurt. The next step in observational equipment came at the time of the experiments by Byard and Kissell. They used a two-stage image converter system with an S-1 photocathode, which was cooled to 0°C to minimize the dark current in the system. Eddy et al. (1973), report that an RCA two-stage S-1 image tube system has a gain of 160 relative to hypersensitized Z-emulsion and an effective gain advantage of about 10 over a single stage S-1 image tube. It seemed in 1972 that observations of the Fe XilI infrared lines could yield considerable information about the physical conditions in the corona, e.g., electron density, elemental abundance, magnetic field orientation, etc., and that better observations of the lines were needed. Three weeks before the 10 July 1972 solar eclipse, one of us (J.M.P.) was offered a new instrument being developed by Tektronix, Inc. to make observations of these lines--a rapid-scanning silicon vidicon spectrometer with both high sensitivity and versatility. This instrument, which had reasonable sensitivity at the infrared wavelengths of the [Fe xm] lines (on the order of 18% of its maximum near-infrared sensitivity), was capable of recording sweeps of a spectral region on magnetic tape in analogue form for later playback and analysis. Its quantum efficiency at 10.800 A is 79/00. We were provided with an engineering model of the spectrometer system (the Tektronix J20/7J20 system), which was used at the eclipse. Coronal lines were detected at this eclipse through cloud cover even with the less-than-adequate tracking system that could be fabricated in the very short time we had before the eclipse. We next used the spectrometer, with a more elaborate mounting system and provision for wavelength calibration at the 30 June 1973 eclipse, observing from the U.S. National Science Foundation eclipse site at Loiengalani, Kenya. Four minutes of observations of the 280 A spectral region centered between 10,747 A and 10,798 A were made in a range of heights between 1.1 and 2.0 Ro.
Infrared coronal lines
343
THEORY
The [Fe xm] transitions under study, 3pj ~ 3po (10,747 A) and 3P 2 ~ 3P l (10,798 A), involve an atom in the 3S 2 3P 2 configuration. Another line in the visible [Fe xnI] spectrum lies at 3388 A. The radiation is governed by magnetic dipole selection rules. A term diagram appears in Fig. 1. ISO TERM DIAGRAM FOR 3t02 CONFIGURATION
OF
,
Fe X l l t
2536 5.53
1279 981.40
,1451
t ~.sz
~
~°z
3388 85.63
r..,!OO0 .0049
5397 po747j, .oo7o7~, ~,3.79 / 3% FIG. 1. Term diagram for the 3s23p 2 ground configuration of Fe xm.
Let us consider the excitation problem and the density dependence of the ratio of the infrared Fe xni lines, similarly to the reasoning of Zirin (1966). First, consider the low density case. Every electron excited to ap: should spontaneously emit the 10,798 A line and then the 10,747 A line, giving a ratio close to the ratio of the Einstein A coefficients, which are approximately equal. The 3P 1 level can be independently excited, which would give an I(10,747 A)/I(10,798 A) ratio on the order to twice the ratio of the Einstein coefficients. It requires more energy for the 15o and tD: levels to be excited than the average photospheric photon has. I D 2 c a n be strongly excited by collisions, so becomes heavily populated only in a high density case. In a high density situation, the number populations are now in accordance with the Boltzmann distribution: t'ln __ gn e - h v / k T
nm
gm
where g, a n d g~ are the statistical weights of the upper and lower states, respectively, hv is the energy difference, T is the temperature and k is Boltzmann's constant. For a temperature on the order of 106 K, e -h~/kr is approximately equal to 1 for all levels. Since g, = 2J + 1, the relative intensity is given by: I(J~J
- ]) = (2J + 1) A2,
where A21 is the Einstein A coefficient for spontaneous emission. This yields an I(10,747
344
JAY M. PASACHOFF and DANIEL F. MUZYKA
A)/I(10.798 A) of 0.85 to 1. Note that this is a limiting case and not necessarily a true representation of the physical picture. The situation where 1(10,747 ,&)/I(10,798 A) is less than 1 has not been observed. Firor and Zirin (1962) discussed the excitational processes for various forbidden transitions in the corona that they observed and concluded that "the normal coronal line is produced by collisional excitation by electron impact and subsequent radiation of the line in question." In most cases the transitions require more than a few eV. They made theoretical calculations of the possible line ratios, including collisional excitation and de-excitation, photoexcitation and de-excitation, and stimulated emission. They neglected ionization and recombination because the probability was small. On the average, their calculations predicted 10,798 to be half as intense as 10,747. Malville (1967) considered the excitation of ground state terms by excitation of levels within 50 eV of the ground state and made calculations for an eighteen level model ion. He also suggested that the fact that the decrease in the equivalent width of 10,747 with height above the limb begins to level off near 1.2 solar radii, which shows that 3P 1 is switching from being collisional to being radiative excited near this height. Malville also summarized the polarization vs. density, relationship for both of the Fe xm infrared lines. Especially in the case of the 10,747 line, the polarization varies steeply with the density. Zirin's (1970) analysis of the same 1962 coronagraph observations of the coronal condensation discussed by Malville (1967) led him to conclude that there is a steep rise in the 10,798 intensity in the condensation when the density is higher. He observed that the line ratio varied from about 7 for "densities less than 10 9 cm-3- and saturates at a value of 2 for densities above that amount. We use the excitation model for the infrared Fe xIII lines of Chevalier and Lambert (1969), slightly adjusted by Finn and Landman (1973). Chevalier and Lambert review the relation of radiative excitation and de-excitation, collisional excitation by protons and collisional excitation via the excited configurations. The first thing they consider is radiative excitation and de-excitation. The formula for radiative excitation is Alu = W g~
A,a
gl exp (hv/k T) - l ' where gu and g~ are the respective statistical weights and W is the dilution factor given by w = ½(1 - , j l
-
1/r2),
where r = height of observation with respect to the center of the disc. W ranges from 0.0 < W < 0.5 for ~ > r > 1. Chevalier and Lambert considered stimulated emission, which is given by A,~ = Alu g~. g~ but found it to be unimportant because it never exceeds 5% of the spontaneous emission rate. The Einstein A coefficients for the two transitions in question are
3P 1 ~ 3Po 3P 2 ~ 3P 1
10,747 10,798
Aut
A,t
13.79 9.66
2.61 1.03
(both S- 1)
These values are calculated for a dilution factor of 0.5 and a temperature of 2 K.
x 10 6
Infrared coronal lines
345
The formula for electron excitation given by Chevalier and Lambert is
St, -
8.63 x 10 -6 ~lu T1/2 oh + 1 e x p ( - h v / k T ) N , , ,
where fit, = collision strength, and collisional de-excitation is given by S,l -
8.63 x 10 -6 ~lu Ta/2 2J, + 1 N,,.
Bahcall and Wolf (1968) suggested that proton collisional excitation might be important in the corona since we can assume that Np (the number of protons) is approximately equal to N,, because of the total ionization of the hydrogen, the most abundant element in the million-degree corona. Their calculation yield limiting excitation probabilities at high temperatures, which they defined as temperatures greater than 1.4 x 107 K. These limiting probabilities are ap o - ~ 3p2~ 3p2j
St, = 3.3 x 1 0 - 9 N i o
3 p 0 --o 3 p 1
Stu = 1.2 x 10 - 9 N v
3p I ~
for each of the ground state transitions. Both Bahcall and Wolf, and Chevalier and Lambert, who used proton collisions as the third component of their model, considered proton excitation of levels above the split ground state levels of Fe xni to be unimportant. At "low" temperatures (T < 1.4 x 106 K), the expression for the proton collisional excitation probability was given by Slu = C 1 exp(C2/T63/5)Np,
where T6 = T/106. C1 and C2 are constants that have been calculated for the ground state transitions. The 3P0----. 3P 1 transition is forbidden in first order theory and has to occur in two steps: 3P 0 ---* 3 P 2 ~ 3P 1 Chevalier and Lambert showed that in a high density situation, i.e., near the limb, this process has considerable effect on the intensity ratio of the two infrared lines of [Fe XllI]. The final component of Chevalier and Lambert's excitation model for the 3s23p2 ground state of Fe xnI was collisional excitation via higher levels. Malville discussed the populating of ground state levels by electron collisional excitation to higher configurations and then radiative de-excitation to levels in the ground state configuration. The collisional excitation rates that Chevalier and Lambert used are given by Sl, = 1.7 x lO-aftffT ~ IO-~°4°W/Tp(W/kT)N,, where W = excitation energy, ftu = absorption oscillator strength, and P ( W / k T ) = a function defined by Van Regemorter (1962). Chevalier and Lambert considered configurations of 3s23p2, 3s23p3d, and 3s23p4s. One example to consider is the collisional excitation of 3s23p2 3 P 2 from the ground state 3s23p2 3 P o. The transition probability of excitation to the 3s3p 3 3S1 excited level is
$1, = 2 x 10 -9 N,,. An excitation to this level is followed, 50~ of the time, by a radiative de-excitation to 3s23p2 3p2, thereby populating one of the higher levels of the split ground state from the lowest level by indirect means. The probability for the net transfer from 3p 2 SPo to 3p 2 3P2 is
Sl, = 9 x l O - l ° N,,. (All these rates are for temperatures of 2 x 106 K.) Other levels are less important to consider. For example, removal of the 3p3d excited configuration from the model had an effect on the I(10,747 /k)/I(10,798 /~) ratio of
346
JAY M. PASACHOFFand DANIEL F. M U Z Y K A Table 1. The [Fe Xln] line ratio for T = 2 x 10 6 K , Ne = 3 x l0 8 cm 3 and W = 0.4. Model description
A. B. C. D.
1(10,747 A)/I(10,798 A)
Radiative and electron collision Model A plus proton collisions Model A plus other excited configuration All four excitation methods
5.66 2.49 3.62 2.37
less than 5~o. The important fact to recall is that electron collisional excitation to different excited configurations has the net effect of populating 3sZ3p 2 3P t and 3sg3p z 3P 2 at the expense of 3sg3p z 3Po. Using a sunspot-minimum model for coronal electron density, Chevalier and Lambert went on to compare various combinations of the four above excitation methods. The following is a table of their results for the T = 2 x 10 6 K , N e = 3 x 108 cm -3 and W = 0.4 case, which they used for comparison with observed intensity ratios at various electron densities (Table 1). Their comparison of the observed intensity ratio of the two lines with an electron density computed from measurements of the K corona showed that model D is the best for this height (r = 1.02 Ro), and, indeed, for all heights. One can readily see from Chevalier and Lambert's results that the ratio is insensitive to temperature in the corona, and is therefore a good probe of the electron density. Chevalier and Lambert's model of the ratio vs. electron density agrees to within 10~ of various observational values. Finn and Landman (1973) added two factors to the model of Chevalier and Lambert: (1) collisional ionization and continuum radiative recombination calculations, and (2) absorption of light from a photosphere with wavelength-dependent limb-darkening characteristics, resulting in a" wavelength-dependent dilution factor. Figure 2 is a graph of the electron density vs. the intensity ratio of the infrared lines, using the calculations of Finn and Landman. The graph shows the dramatic association between the intensity ratio and N,,. EQUIPMENT
The Tektronix spectrometer used for our eclipse observations was developed for situations where one has to be able to record spectral information, such as the intensity 1.6
40 30 T=2xl0
1.4
6
K
20 1.2 O~
1.0
,o
0.8
8 7 6 5
0.6
4
H
~,¢ 0 ,< H
t.-
o__ H
_d I-4
3 0.4 2
0.2
0
P 6
7
8
9
TO
FIG. 2. Log N e vs. I(10,747 A)/1(10,798 A), based on the calculations of Chevalier and Lambert (1969) as adjusted by Finn and Landman (1973).
Infrared coronal lines
ENTRANCE SLIT HORIZ , , ~
"
"-
347
~
- _ _
.
O$CILLO6COPE
,,.'"
J ;I,'
/
i/
ELEC I...J"ECTRAL1 NORMAL-|
PROCESSlNC1-~
IZlNG J
VERT
SCANNINGRAMP
I
$ TR,GGER,NG
]
FIG. 3. Block diagram of the spectrometer unit.
of radiation that enters the system, that varies rapidly in time or when it is necessary to vary rapidly the wavelength range being observed without extensive instrumental recalibration. At the eclipse, we took advantage of both of these capabilities and the sensitivity of the vidicon in the near infrared. We made height scans of the intensity of 10,747 and 10,798 and, for a very short time, searched for new lines in the spectral range from 8000 A to 11,000 A. The basic configuration of the instrument is shown in Fig. 3. We had pre-production model A-01 at the eclipse. Parameters of the instrument configuration during the eclipse can be seen in Table 2. Table 2. RSS settings at eclipse Grating: Slit: Time/scan: Center of display: Display width:
"B" (blazed in the infrared) 100 pm 20 ms 10800 A (instrumental) 278 A (digitized channels 1 to 256)
The coronal image was reflected by a 40 cm coelostat into a 20 cm f/10 Celestron telescope and focused on the slit of the spectrometer. Beyond this, the light entered the self-contained system. In the system, it is dispersed by a Czerny-Turner monochromator, which images the spectrum onto the vidicon target. The vidicon consists of a reverse-based photodiode array grown on a silicon wafer. Light incident upon the target produces electron-hole pairs, with the hole migrating to the side of the target, which is addressed by the electron beam. The electron is conducted to ground. The holes accumulate until they are scanned off; the current
348
JAY M. PASACHOFF and DANIEL F. MUZYKA
needed to recharge the diodes is proportional to the incident light intensity. The fact that the spectrum is actually imaged on the vidicon allows the spectrometer to gather data over a wide range of wavelengths. Integration can be achieved internally for 50, 100, 200, 500, or 1000 ms, representing the scan time plus the possible retardations of the scanning electron beam between readings. After the array is scanned, the spectrum is displayed on a Tektronix 7000-series oscilloscope in which a special plug-in module contains half of the spectrometer electronics. The instrument has two gratings. Grating A gives a 4000 A bandpass and is variable in 1000 A steps, while grating B is continuously variable in central wavelength and displays approximately 400 A at 5000 A, and somewhat less with the infrared. A special non-standard grating blazed in the infrared was installed for our eclipse observations. The grating was blazed in the infrared at 10,000 /~ and has about 63~o efficiency at 10,800 A. It was subsequently determined that the actual spectral range displayed during the eclipse was 278/~. Though the spectral sensitivity of the instrument at 10,800 A is only 18~o of the maximum spectral sensitivity in the range of 7000 A to 11,000 A, it is still much better than other methods. The sensitivity corresponds to the convolution of the vidicon quantum efficiency of 7 ~ and the response of the electronics. Its signal at 10,800 A is 3.5 amp A/w. It was of much concern to us that the noise figure of the instrument doubles every 10°C. We had hoped to receive a spectrometer unit that had a liquid/gaseous nitrogen cooling system for the detector, but the technical revisions of the original instrument were not completed by the time of the eclipse and temperature control in Africa was difficult. The importance of the noise figure can more easily be shown by a simple development of the equations involved. The (spot) noise figure F (IRE, 1960) of a system is defined as F _ W.. + W. Wo" ' where W, is the output power attributable to the system and W0, is the output power attributable to the input source. (The gain of a system is defined as W0,/source input power.) It can be shown that this reduces to
F = I + T/ To or
T = (F - 1)To
where T = system noise temperature; To = ambient temperature. The sensitivity or minimum detectable temperature and therefore the minimum detectable flux is given by the rms noise temperature of a system Tsys
Tmi, oC , ~ ,
where T~ysis the system temperature from the previous formula, and t is the integration time. Let us consider an example. If F = 3, To = 290 K, Zy~ by the given formula is 580 K. If one were to integrate for the same length of time and increased the ambient temperature by 10°C to To = 300 K, and therefore double the noise figure to F = 6, the system temperature increases to Zys = 1500 K. This indicates a decrease by a factor of 2.6 in the minimum detectable flux of the instrument. Obviously, this is not a rigorous development of this subject but it does point out our cause for concern at the eclipse. It should be noted that a doubling of a noise figure of approximately 1.5 (To = 290 D), a more realistic figure than F = 3, would yield a decrease in the minimum detectable flux on the order of 4.
Infrared coronal lines DATA
349
COLLECTION
The Williams group was pooled for logistical purposes with other university and research institution groups that were given grants from the National Science Foundation. The National Science Foundation chose to make its prime site in Kenya, where there would be 5 minutes of totality, instead of sites in West Africa, where totality was longer, because of the expected presence of dust in the air in West Africa as a result of the extended drought there. The particular site at the Kenyan oasis of Loiengalani was chosen because it provided a view of the eclipse over water, which would tend to stabilize the images. As it turned out, a dust storm produced about 90~o obscuration on the day of the eclipse at the West African site in Mauritania and ruined many experiments. But the Kenyan site, where we were located, had mostly clear skies. Some small clouds in the sky at totality cleared the sun only 45 seconds before second contact (for a more thorough report of eclipse site activity, see Pasachoff, 1974). Most of our equipment had been shipped to the site at Loiengalani, Kenya; however, we brought the monochromator/vidicon unit of the instrument on the aircraft with us. It was kindly flown from Nairobi to Loiengalani by the Royal Kenya Air Force. The first order of business at the site after arrival was to construct a building to house the spectrometer. As the winds were approximately 60 mph each night, it was useful to construct a solid structure to house the equipment. A floor plan of the building appears as Fig. 4. We also constructed several concrete piers both inside and outside the building to use as stable platforms for the spectrometer and tracking equipment. Unfortunately, because of funding limitations, all the equipment arrived at the last moment before shipping, and was assembled together for the first time on site in Kenya. The effect of the wind-blown sand can be recalled from the example of the air conditioner that we brought with us. Dust got into the contacts of one of the relays and
I 1 3 4 5 6 7 8 9 l0 ii
OSCILLOSCOPE C A L I B R A T I O N LAMPS SPECTROMETER CELESTRON OSCILLOSCOPE V O I C E TAPE R E C O R D E R N O T I O N PICTURE C A M E R A COELOSTAT I N S T R U M E N T A L TAPE RECORDER REFRIGERATOR AIR CONDITIONER
10
WILLIAMS COLLEGE
K-17
FIG. 4. Floor plan of the Williams College spectrometer building at the N.S.F. eclipse site at Loiengalani, Kenya. The equipment in the shed included a film refrigerator (10), two oscilloscopes (1 and 5), the air conditioner, built into the wall (11), the Honeywell instrumental tape recorder (8), the spectrometer unit (3), a box with several calibration lamps (including neon, mercury and argon sources) (7), a Mitchell motion picture camera, used to record the actual pointing of the coelostat (8), a tape recorder for recording operating commands made by the group (6), and the 20 cm Celestron telescope (4). J,V.A. 19/4
350
JAY M. PASACHOFF and DANIEL F. MUZYKA
caused it to short out. There is no room here to go into the story of how a slightly smaller 110 V replacement was found and delivered to us in northern Kenya within 24 hours. A major problem at the site was temperature control. Temperatures typically reached 43°C (ll0°F) during the afternoon. Since an internally-cooled spectrometer was not available at the time of the eclipse, we had to try other means of keeping the temperature of the vidicon at the lowest possible level during the eclipse. Besides the air conditioner, we attempted to cool the unit further during the eclipse by placing the remnants of a shipment of dry ice around the unit at the time of the eclipse. (The dry ice had been flown in from Nairobi for the use of other observers.) All of these measures were helpful, but during the actual eclipse a large Lexan window we had installed on the east side of the building had to be removed to eliminate its absorption of incoming infrared radiation. This meant that the building was partially open to the surrounding atmosphere and, despite the cooling efforts, the temperature in the room was approximately 27°C. The temperature in the spectrometer housing itself was probably higher still. This led to a diminished sensitivity for the instrument, as previously discussed. During the eclipse, some of the light from the coelostat was reflected into the Mitchell movie camera. This camera recorded the pointing of the coelostat and therefore the pointing of the telescope. Another diagonal mirror reflected some of the light back to the coelostat mount, to a Nikon camera with cross-hairs, through which the coelostat operator could monitor his adjustments of the pointing. Another camera mounted on top of the Celestron and an adjacent bore-sighted telescope allowed the spectrometer operators inside the house to monitor the tracking as well. Scans of live emission from several spectral lamps (Pen-Ray variety) were taken at intervals before, during and after the eclipse. The lamps were mounted in a box with a slit. Radiation from the box could be reflected into the spectrometer in place of the radiation from the corona by flipping a mirror set between the spectrometer and the Celestron. A filter wheel was placed just before the entrance slit to the spectrometer. Polaroids were mounted in the wheel at 3 polarization angles and one slot was left blank to pass unpolarized light to the spectrometer. The output of the spectrometer unit went to the oscilloscope marked 5 in the floor plan, where the coupling module for the spectrometer was located. The spectra were displayed dynamically on both oscilloscopes. At the eclipse, we recorded the data in FM mode on several channels of a Honeywell 3600C instrumental tape recorder and recorded trigger/timing pulses on other channels. For most of the 5 minutes of totality, we observed in the spectral region that included 10,747 and 10,798. Most of these data were taken of unpolarized radiation. We devoted less than 1 minute at the end of the eclipse to scanning through the spectrum from 7000 h to 11,000 A. Time limitations prevented our observing the 3388 ,~ line of Fe xIn. ANALYSIS
When the data were returned to the U.S., they were first digitized for computer processing. This was carried out at the Tektronix headquarters in Beavertom Oregon. This first digitization referred to as DIG-I, was carried out by playing the analogue tape from the Honeywell instrumental tape recorder into a Tektronix Digital Processing Oscilloscope and recording the data with a PDP 11/25. The intensity level of each scan was sampled at 256 equally spaced points. Each 128 scans were averaged (each scan representing 0.02 s of integration) into "blocks". For the time of the eclipse alone, if the individual scans were left unaveraged, we would have 4 x 10 6 pieces of information. By averaging the data into blocks of 128 scans, the number of points was reduced to approximately 3 x 104. We later discovered that problems arose in the triggering during the first digitization. Because of these problems, some scans and therefore data were lost. Therefore a second digitization of the data was carried out, which we refer to as DIG-2. Each block of
Infrared coronal lines
351
DIG-2 data represented about 5 s of actual observation time. By February of 1975 most of the digitization problems were corrected and the digital information was transferred from DEC compatible DECTAPE to standard 9 track, 800 bpi magnetic tape and punched cards. We deduced the spectral resolution of the instrument by degrading a section of the Kitt Peak Solar Atlas. We degraded the spectrum between 10,000 A to 10,800 A by moving averages from 0.02 A resolution to resolutions of 1 A to 6 A in order to match the resolution on the vidicon spectra. At the same time, we had available a "solar atlas" made with the eclipse unit. A comparison of the two sets of photospheric spectra showed that the eclipse A-01 unit produced spectra in the region of the two Fe xIn infrared lines of 2 A resolution. Another set of values that had to be determined was actual centers of each spectrum, and what the dispersion of the spectra were. The instrument was centered on 10,800 A, but there is an uncertainty of 10 A in the position of the intensified spot, called the "marker", on the oscilloscope readout. The actual centers of the spectra recorded at the eclipse were determined from calibration source lines. The dispersion corresponded to a total display of 278 A at our infrared wavelength. [re xm] to74rl 3 0 June 1973 F.nllpse at Lolengolnnlo Kenya i
60
h " 1.4 Rn
30
-30
-60 0
j I00
i 200
FIG. 5. An averaged spectrum showing the 10,747 A line of Fe Xlll at a height of 1.4 Ro from the data taken during the 30 June 1973 total solar eclipse.
The spectra were analyzed on the Xerox 530 computing system at Williams College. The upward drifting of successive spectra was corrected by the fitting and subtraction of a first degree polynomial from all of the spectra. Edge effects were dealt with by eliminating the end channels in all arithmetic operations carried out on the spectra. A series of spectra were found to be a good reference baseline, were averaged and then were subtracted from the already "normalized" spectra. An emission feature detected at 10,747 A on the DIG-2 d a t a corresponding to h = 1.4 Ro is shown in Fig. 5. Despite further averaging no feature was detected at 10,798 A and so at this time it is only possible to state a lower limit on the I(10,747 A)/I(10,798 A) ratio of 5/1. Using the Chevalier and Lambert model as adjusted by Finn and Landman for a temperature of T = 2 × 106 K (Fig. 2), this corresponds to an upper limit for the electron density of 3 × 107 for h = 1.4 Ro (Pasachoff and Muzyka, 1975). Flower and Nussbaumer (1974) have recently derived mean values for the electron density corresponding to height of 1.02 Ro from observations of permitted Fe xm lines. The lines in the 200 to 300 • range were observed by Malinovsky and Heroux (1973) from a rocket launched in 1969. Allowing for the difference in heights of observation, the values from the infrared and ultraviolet lines are consistent with each other. The limit for the ratio of intensities we have set gives better agreement between electron densities derived from the Fe xni line ratios and electron densities derived
352
JAY M. PASACHOFFand DANIEL F. MUZYKA
from non-Fe xIII methods than did the previous eclipse measurements of Kurt, which Chevalier and Lambert suggested were affected by scattered light. Our results support this conclusion. We are continuing the data reduction by redigitizing the data with a new PDP-11/10 at Williams. We hope that this further data reduction will sufficiently improve the signalto-noise ratio to detect the 10,798 line at this and other coronal heights, and thus give us more values of the electron density.
ACKNOWLEDGEMENTS We thank Mr. J. Philip Schierer and Tektronix, Inc. for lending us the spectrometer and for giving various other types of assistance in several stages of this project. Dr. William Weir of Reed College and Tektronix has kindly assisted with programming for setting up the Digital Processing Oscilloscope at Tektronix. Messrs. Stuart N. Vogel '75 and Dan Stinebring '75 also participated in the expedition, and were a vital part of all phases of this experiment in Africa. Mr. Mark Chaffee '77 made the measurements of the pointing. We also thank many other individuals at Williams College for their help. Prof. Stuart J. B. Crampton was of valuable aid in transferring digital data for us. Prof. Howard Stabler designed an FM modulator/ demodulator, which was built by Mr. Emile Ouellette. Messrs. Joseph McCann and Albert St. Pierre expertly constructed mountings, packing crates, and other necessities. We thank Professor Donald H. Menzel of the Center for Astrophysics for useful conversations and for introducing us to the techniques of eclipse observing that he has perfected over the years. Donald A. Cooke '75 assisted with the equipment for this experiment in Williamstown, and then went on to a site in West Africa with Professor Menzel to carry out another spectrographic experiment. We are grateful to him for taking on this responsibility. We are especially thankful to Mr. Ronald R. La Count, U.S. Solar Eclipse Coordinator at the National Science Foundation, Dr. G. William Curtis, Director of the N.S.F. expedition, and Mr. Eugene Prantner and the entire support staff of the Expedition from the National Center for Atmospheric Research in Boulder. We also thank Dr. Harold Lane of the N.S.F. for his encouragement of this project. We are grateful to Mr. Bud Corbin of the Honeywell Corporation for making their 5600C instrumental tape recorder available to us both in Africa and in Oregon. We also acknowledge with thanks the assistance of Mark Armstead, Inc. of Hollywood, California, for lending us the Mitchell movie cameras used for some of the time-lapse movies, and of Messrs. Dennis Kane and Sidney Platt of the National Geographic Society for lending us an Arriflex movie camera for other time-lapse work. We wish to thank Mrs. Mary G. Smith, and others at the National Geographic Society for their encouragement over and above the financial support that we have acknowledged below. We also appreciate the assistance of the Photographic Division of National Geographic, especially Messrs. Robert Gilka and John Fletcher and the people in the black and white and color labs who worked very hard on technically tricky negatives, for providing camera equipment and for handling all of the film processing and printing. We thank Dr. A. Keith Pierce and Mr. Larry Testerman of the Kitt Peak National Observatory for providing us with their solar Atlas in advance of publication. Dr. Robert F. Howard of the Hale Observatories and Dr. Stephen A. Schoolman, then of Caltech and now of the Lockheed Solar Observatory, were of great assistance in the definition stage of this experiment at the time the spectrometer first became available to us. We would like to thank the Computer Center at Williams for providing D.F.M. with time off from his position and for providing large amounts of time on the computer. We were supported in part by National Science Foundation grants GP37187 and MPS 74-23531, and grants from the National Geographic Society and the Research Corporation. We would also like to acknowledge support from a grant by the Alfred P. Sloan Foundation to Williams College, and from the President's Discretionary Fund at Williams. The Government of Kenya deserves many thanks for the hospitality and services it provided for us and for the entire expedition.
REFERENCES Bahcall, J. and Wolf, R. A. (1968) Ann. Astrophys. 29, 343. Byard, P. and Kissell, K. (1971) Solar Phys. 21, 351. Chevalier, R. and Lambert, D. (1969) Solar Phys. 10, 115. Dumont, J.-P. and Perche, J.-C. (1964) 12me Coll. Int. Astrophys. Libge p. 86. Eddy, J. A., Firor, J. W. and Lee, R. H. Bull. Am. Astron. Soc. 72. Eddy, J. A., Lee, R. and Emerson, J. (1973) Solar Phys. 30, 351. Eddy, J. A. and Malville J. (1967) Astrophys. J. 150, 289. Edl6n, B. (1942) Z. Astrophys. 22, 30. Finn, G. D. and Landman, D. (1973) Solar Phys. 30, 381. Firor, J. and Zirin, H. (1962) Astrophys. J. 135, 122. Fisher, R. and Pope, T. (1971) Solar Phys. 20, 389. Flower, D. R. and Nussbaumer, H. (1974) Astron. Astrophys. 31, 353. IRE. Standards on Receivers, Definition of Terms, Proc. IRE 40, 1681. Kurt, V. G. (1962) Soviet Astron. 6, 349.
Infrared coronal lines
353
Kurt, V. G. (1963) Soviet Astron. 6, 620. Lyot, B. (1939) M.N.R.A.S. 99, 580. Malinovsky, M. and Heroux, L. (1973) Astrophys. J. 181, 1009. Malville, J. (1967) Astrophys. J. 148, 299. Pasachoff, J, M. (1974) unpublished report. Pasachoff, J. M. and Muzyka, D. F. (1975) Bull. Am. Astron. Soc. 7, 409, Ratier, G. and Rozelot, J.-P. (1971) Solar Phys. 23, 394. Shklovskii, J. S. (1965) Physics of the Solar Corona, Pergamon Press, Oxford. Van Regemorter, H. (1962) Astrophys. J. 136, 906. Wl6rick, G., Dumont, J.-P. and Perche, J.-C. (1963) in The Solar Corona, Ed,, J. W. Evans, p. 177, Academic Press, New York. Zirin, H. (1966) The Solar Atmosphere, Blaisdell. Zirin, H. (1970) Solar Phys. 11, 497.
Vol. 19, pp. 341-353.PergamonPress, 1976.Printedin Great Britain
INFRARED CORONAL LINES (Observations of Infrared I-Iron XIII] at the 1973 Total Solar Eclipse) JAY M . PASACHOFF a n d DANIEL F. MUZYKA Williams College--Hopkins Observatory, Williamstown, Massachusetts.
IT HAS only been recently that vidicon systems have gained sufficient sensitivity as to allow their use in ground-based infrared studies of the solar corona. This paper reports on the use of a silicon vidicon spectrometer for ground-based eclipse studies of infrared coronal lines. Two forbidden spectral emission lines in the ground state of Fe xm in the near infrared at 10,747 A (3P 1 ---. 3po) and 10,798 ]k (aPE --~ aP1) are the strongest coronal lines in line-to-continuum ratio accessible to ground-based observers. It has been shown that the ratio of the intensity of these lines is insensitive to temperature and provides a good measure of the electron density in corona since one level is radiatively and the other collisionally populated. Simultaneous observations of the 10,747 A and 10,798 A lines therefore provide a sensitive probe of coronal electron density. Both coronagraphic and eclipse observations of these lines have been made but all observations have been severely handicapped by the fact that the lines lie beyond the effective limit of most photographic film and on the upper end of the wavelength limit of the S-1 photocathode of the most commonly employed image tube. At the total solar eclipse of June 30, 1973 in Kenya, we observed the spectral region that included the two lines with a Tektronix J20/TJ20 rapid scanning silicon vidicon spectrometer. We scanned the corona in height from 1.1 to 2.0 RQ. We report here on the observations and their reduction, as well as describe the instrumentation and the theory of the [Fe xIII] lines. INTRODUCTION
Using the coronograph at the Pic du Midi observatory in 1936, Bernard Lyot made a survey of the spectrum of the corona from 3550 to 12,000 A. Conspicuous among the emission lines he found in the coronal spectra were two at 10,747 and 10,798 A that have the two highest line-to-continuum ratios of any of the lines discovered. They showed up "after an extensive trace of the spectrum without coronal radiation" (Lyot, 1939). Edl~n (1942) identified the lines 3 years later as the aP x --~ 3P o (10,747 A) and aP 2 ~ 3P 1 (10,798 A) forbidden transitions of the 3sZ3P2, ground state of Fe xIn, twelvetimes-ionized iron. The importance of these lines was discussed by Firor and Zirin (1962), who showed that the 10,798 A line could not possibly be excited by photoexcitation. This realization and further development of the excitation theory has led to the understanding that one level is basically radiatively and the other collisionally populated, and that a ratio of the line intensities is an accurate tool for measuring electron density in the corona. The theory has been elaborated by Malville (1967), Chevalier and Lambert (1969), Ratier and Rozelot (1972), and Finn and Landman (1973). Though minor corrections in various parameters have been suggested, the theory as set forth by Chevalier and Lambert is now widely accepted. It has been shown both theoretically and observationally that the intensity ratio I(I0,747 ]~)/I(10,798 ]~) is on the order of 7 or less for electron densities much less than 109 c m - 3 and saturates at a value of 2 for densities above that amount. The modern era of observations of these lines began in 1936, when Soviet observers at Pulkovo and Kislovodsk observed the lines with an image tube system (as reported 341
342
JAY M. PASACHOFFand DANIEL F. MUZYKA
by Shkiovsky, 1965). Despite the discovery and identification of the lines in the 1930's, it was 20 years before anyone again observed them. The major problem with observations in those 20 years was that observations could not have been made with any more resolution or structure showing than those made by Lyot--the film limitation on near infrared photography continued. Eclipse observations such as those we and others have since made would have been even more difficult to carry out had the only method of observing been with photographic film. Non-eclipse, coronograph, observations resumed at Climax when observations of 10,747 and 10,798 were made in 1959, subsequently studied by Firor and Zirin (1962), and in 1962 on dates that closely coincided with the New Guinea total eclipse, subsequently studied by Malville (1967) and Zirin (1970). Other non-eclipse observations were reported by W16rick et al. (1963), Dumont and Perche (1964), Fisher and Pope (1971) and Ratier and Rozelot (1972). During the total solar eclipse of 12 November 1966, Byard and Kissell (1971) observed the [Fe xm] lines from the NASA 990 off the southern coast of Brazil. Overall, the observations yield electron densities in the range of l0 s to 109 cm -3, depending upon the height above the limb at which the observations were taken, and on the values of excitation parameters used in the calculations. Other observations of 10,747 A were made to measure line polarization at the 1965 eclipse (Eddy and Malville, 1967) and at the 1966 eclipse (Eddy et al., 1967, 1973). A major problem, as pointed out earlier, is that all observations have been hampered by the low sensitivity of the observing instruments and film. Lyot used Eastman Z-emulsion hypersensitized with ammonia in his spectral line search in the corona. Kurt observed the lines with an electron-optical image converter with an oxygen-cesium cathode. This system had a reported sensitivity of 100 times over the system used by Lyot. A single-stage converter was also used for the Climax observations. The tube, an RCA 6032, had an S-1 photocathode, which has a quantum efficiency at 10,800 A of approximately 0.4%. The converter's output was focused on a phosphor screen to be photographed or viewed. This unit was perhaps slightly more efficient than the unit used by Kurt. The next step in observational equipment came at the time of the experiments by Byard and Kissell. They used a two-stage image converter system with an S-1 photocathode, which was cooled to 0°C to minimize the dark current in the system. Eddy et al. (1973), report that an RCA two-stage S-1 image tube system has a gain of 160 relative to hypersensitized Z-emulsion and an effective gain advantage of about 10 over a single stage S-1 image tube. It seemed in 1972 that observations of the Fe XilI infrared lines could yield considerable information about the physical conditions in the corona, e.g., electron density, elemental abundance, magnetic field orientation, etc., and that better observations of the lines were needed. Three weeks before the 10 July 1972 solar eclipse, one of us (J.M.P.) was offered a new instrument being developed by Tektronix, Inc. to make observations of these lines--a rapid-scanning silicon vidicon spectrometer with both high sensitivity and versatility. This instrument, which had reasonable sensitivity at the infrared wavelengths of the [Fe xm] lines (on the order of 18% of its maximum near-infrared sensitivity), was capable of recording sweeps of a spectral region on magnetic tape in analogue form for later playback and analysis. Its quantum efficiency at 10.800 A is 79/00. We were provided with an engineering model of the spectrometer system (the Tektronix J20/7J20 system), which was used at the eclipse. Coronal lines were detected at this eclipse through cloud cover even with the less-than-adequate tracking system that could be fabricated in the very short time we had before the eclipse. We next used the spectrometer, with a more elaborate mounting system and provision for wavelength calibration at the 30 June 1973 eclipse, observing from the U.S. National Science Foundation eclipse site at Loiengalani, Kenya. Four minutes of observations of the 280 A spectral region centered between 10,747 A and 10,798 A were made in a range of heights between 1.1 and 2.0 Ro.
Infrared coronal lines
343
THEORY
The [Fe xm] transitions under study, 3pj ~ 3po (10,747 A) and 3P 2 ~ 3P l (10,798 A), involve an atom in the 3S 2 3P 2 configuration. Another line in the visible [Fe xnI] spectrum lies at 3388 A. The radiation is governed by magnetic dipole selection rules. A term diagram appears in Fig. 1. ISO TERM DIAGRAM FOR 3t02 CONFIGURATION
OF
,
Fe X l l t
2536 5.53
1279 981.40
,1451
t ~.sz
~
~°z
3388 85.63
r..,!OO0 .0049
5397 po747j, .oo7o7~, ~,3.79 / 3% FIG. 1. Term diagram for the 3s23p 2 ground configuration of Fe xm.
Let us consider the excitation problem and the density dependence of the ratio of the infrared Fe xni lines, similarly to the reasoning of Zirin (1966). First, consider the low density case. Every electron excited to ap: should spontaneously emit the 10,798 A line and then the 10,747 A line, giving a ratio close to the ratio of the Einstein A coefficients, which are approximately equal. The 3P 1 level can be independently excited, which would give an I(10,747 A)/I(10,798 A) ratio on the order to twice the ratio of the Einstein coefficients. It requires more energy for the 15o and tD: levels to be excited than the average photospheric photon has. I D 2 c a n be strongly excited by collisions, so becomes heavily populated only in a high density case. In a high density situation, the number populations are now in accordance with the Boltzmann distribution: t'ln __ gn e - h v / k T
nm
gm
where g, a n d g~ are the statistical weights of the upper and lower states, respectively, hv is the energy difference, T is the temperature and k is Boltzmann's constant. For a temperature on the order of 106 K, e -h~/kr is approximately equal to 1 for all levels. Since g, = 2J + 1, the relative intensity is given by: I(J~J
- ]) = (2J + 1) A2,
where A21 is the Einstein A coefficient for spontaneous emission. This yields an I(10,747
344
JAY M. PASACHOFF and DANIEL F. MUZYKA
A)/I(10.798 A) of 0.85 to 1. Note that this is a limiting case and not necessarily a true representation of the physical picture. The situation where 1(10,747 ,&)/I(10,798 A) is less than 1 has not been observed. Firor and Zirin (1962) discussed the excitational processes for various forbidden transitions in the corona that they observed and concluded that "the normal coronal line is produced by collisional excitation by electron impact and subsequent radiation of the line in question." In most cases the transitions require more than a few eV. They made theoretical calculations of the possible line ratios, including collisional excitation and de-excitation, photoexcitation and de-excitation, and stimulated emission. They neglected ionization and recombination because the probability was small. On the average, their calculations predicted 10,798 to be half as intense as 10,747. Malville (1967) considered the excitation of ground state terms by excitation of levels within 50 eV of the ground state and made calculations for an eighteen level model ion. He also suggested that the fact that the decrease in the equivalent width of 10,747 with height above the limb begins to level off near 1.2 solar radii, which shows that 3P 1 is switching from being collisional to being radiative excited near this height. Malville also summarized the polarization vs. density, relationship for both of the Fe xm infrared lines. Especially in the case of the 10,747 line, the polarization varies steeply with the density. Zirin's (1970) analysis of the same 1962 coronagraph observations of the coronal condensation discussed by Malville (1967) led him to conclude that there is a steep rise in the 10,798 intensity in the condensation when the density is higher. He observed that the line ratio varied from about 7 for "densities less than 10 9 cm-3- and saturates at a value of 2 for densities above that amount. We use the excitation model for the infrared Fe xIII lines of Chevalier and Lambert (1969), slightly adjusted by Finn and Landman (1973). Chevalier and Lambert review the relation of radiative excitation and de-excitation, collisional excitation by protons and collisional excitation via the excited configurations. The first thing they consider is radiative excitation and de-excitation. The formula for radiative excitation is Alu = W g~
A,a
gl exp (hv/k T) - l ' where gu and g~ are the respective statistical weights and W is the dilution factor given by w = ½(1 - , j l
-
1/r2),
where r = height of observation with respect to the center of the disc. W ranges from 0.0 < W < 0.5 for ~ > r > 1. Chevalier and Lambert considered stimulated emission, which is given by A,~ = Alu g~. g~ but found it to be unimportant because it never exceeds 5% of the spontaneous emission rate. The Einstein A coefficients for the two transitions in question are
3P 1 ~ 3Po 3P 2 ~ 3P 1
10,747 10,798
Aut
A,t
13.79 9.66
2.61 1.03
(both S- 1)
These values are calculated for a dilution factor of 0.5 and a temperature of 2 K.
x 10 6
Infrared coronal lines
345
The formula for electron excitation given by Chevalier and Lambert is
St, -
8.63 x 10 -6 ~lu T1/2 oh + 1 e x p ( - h v / k T ) N , , ,
where fit, = collision strength, and collisional de-excitation is given by S,l -
8.63 x 10 -6 ~lu Ta/2 2J, + 1 N,,.
Bahcall and Wolf (1968) suggested that proton collisional excitation might be important in the corona since we can assume that Np (the number of protons) is approximately equal to N,, because of the total ionization of the hydrogen, the most abundant element in the million-degree corona. Their calculation yield limiting excitation probabilities at high temperatures, which they defined as temperatures greater than 1.4 x 107 K. These limiting probabilities are ap o - ~ 3p2~ 3p2j
St, = 3.3 x 1 0 - 9 N i o
3 p 0 --o 3 p 1
Stu = 1.2 x 10 - 9 N v
3p I ~
for each of the ground state transitions. Both Bahcall and Wolf, and Chevalier and Lambert, who used proton collisions as the third component of their model, considered proton excitation of levels above the split ground state levels of Fe xni to be unimportant. At "low" temperatures (T < 1.4 x 106 K), the expression for the proton collisional excitation probability was given by Slu = C 1 exp(C2/T63/5)Np,
where T6 = T/106. C1 and C2 are constants that have been calculated for the ground state transitions. The 3P0----. 3P 1 transition is forbidden in first order theory and has to occur in two steps: 3P 0 ---* 3 P 2 ~ 3P 1 Chevalier and Lambert showed that in a high density situation, i.e., near the limb, this process has considerable effect on the intensity ratio of the two infrared lines of [Fe XllI]. The final component of Chevalier and Lambert's excitation model for the 3s23p2 ground state of Fe xnI was collisional excitation via higher levels. Malville discussed the populating of ground state levels by electron collisional excitation to higher configurations and then radiative de-excitation to levels in the ground state configuration. The collisional excitation rates that Chevalier and Lambert used are given by Sl, = 1.7 x lO-aftffT ~ IO-~°4°W/Tp(W/kT)N,, where W = excitation energy, ftu = absorption oscillator strength, and P ( W / k T ) = a function defined by Van Regemorter (1962). Chevalier and Lambert considered configurations of 3s23p2, 3s23p3d, and 3s23p4s. One example to consider is the collisional excitation of 3s23p2 3 P 2 from the ground state 3s23p2 3 P o. The transition probability of excitation to the 3s3p 3 3S1 excited level is
$1, = 2 x 10 -9 N,,. An excitation to this level is followed, 50~ of the time, by a radiative de-excitation to 3s23p2 3p2, thereby populating one of the higher levels of the split ground state from the lowest level by indirect means. The probability for the net transfer from 3p 2 SPo to 3p 2 3P2 is
Sl, = 9 x l O - l ° N,,. (All these rates are for temperatures of 2 x 106 K.) Other levels are less important to consider. For example, removal of the 3p3d excited configuration from the model had an effect on the I(10,747 /k)/I(10,798 /~) ratio of
346
JAY M. PASACHOFFand DANIEL F. M U Z Y K A Table 1. The [Fe Xln] line ratio for T = 2 x 10 6 K , Ne = 3 x l0 8 cm 3 and W = 0.4. Model description
A. B. C. D.
1(10,747 A)/I(10,798 A)
Radiative and electron collision Model A plus proton collisions Model A plus other excited configuration All four excitation methods
5.66 2.49 3.62 2.37
less than 5~o. The important fact to recall is that electron collisional excitation to different excited configurations has the net effect of populating 3sZ3p 2 3P t and 3sg3p z 3P 2 at the expense of 3sg3p z 3Po. Using a sunspot-minimum model for coronal electron density, Chevalier and Lambert went on to compare various combinations of the four above excitation methods. The following is a table of their results for the T = 2 x 10 6 K , N e = 3 x 108 cm -3 and W = 0.4 case, which they used for comparison with observed intensity ratios at various electron densities (Table 1). Their comparison of the observed intensity ratio of the two lines with an electron density computed from measurements of the K corona showed that model D is the best for this height (r = 1.02 Ro), and, indeed, for all heights. One can readily see from Chevalier and Lambert's results that the ratio is insensitive to temperature in the corona, and is therefore a good probe of the electron density. Chevalier and Lambert's model of the ratio vs. electron density agrees to within 10~ of various observational values. Finn and Landman (1973) added two factors to the model of Chevalier and Lambert: (1) collisional ionization and continuum radiative recombination calculations, and (2) absorption of light from a photosphere with wavelength-dependent limb-darkening characteristics, resulting in a" wavelength-dependent dilution factor. Figure 2 is a graph of the electron density vs. the intensity ratio of the infrared lines, using the calculations of Finn and Landman. The graph shows the dramatic association between the intensity ratio and N,,. EQUIPMENT
The Tektronix spectrometer used for our eclipse observations was developed for situations where one has to be able to record spectral information, such as the intensity 1.6
40 30 T=2xl0
1.4
6
K
20 1.2 O~
1.0
,o
0.8
8 7 6 5
0.6
4
H
~,¢ 0 ,< H
t.-
o__ H
_d I-4
3 0.4 2
0.2
0
P 6
7
8
9
TO
FIG. 2. Log N e vs. I(10,747 A)/1(10,798 A), based on the calculations of Chevalier and Lambert (1969) as adjusted by Finn and Landman (1973).
Infrared coronal lines
ENTRANCE SLIT HORIZ , , ~
"
"-
347
~
- _ _
.
O$CILLO6COPE
,,.'"
J ;I,'
/
i/
ELEC I...J"ECTRAL1 NORMAL-|
PROCESSlNC1-~
IZlNG J
VERT
SCANNINGRAMP
I
$ TR,GGER,NG
]
FIG. 3. Block diagram of the spectrometer unit.
of radiation that enters the system, that varies rapidly in time or when it is necessary to vary rapidly the wavelength range being observed without extensive instrumental recalibration. At the eclipse, we took advantage of both of these capabilities and the sensitivity of the vidicon in the near infrared. We made height scans of the intensity of 10,747 and 10,798 and, for a very short time, searched for new lines in the spectral range from 8000 A to 11,000 A. The basic configuration of the instrument is shown in Fig. 3. We had pre-production model A-01 at the eclipse. Parameters of the instrument configuration during the eclipse can be seen in Table 2. Table 2. RSS settings at eclipse Grating: Slit: Time/scan: Center of display: Display width:
"B" (blazed in the infrared) 100 pm 20 ms 10800 A (instrumental) 278 A (digitized channels 1 to 256)
The coronal image was reflected by a 40 cm coelostat into a 20 cm f/10 Celestron telescope and focused on the slit of the spectrometer. Beyond this, the light entered the self-contained system. In the system, it is dispersed by a Czerny-Turner monochromator, which images the spectrum onto the vidicon target. The vidicon consists of a reverse-based photodiode array grown on a silicon wafer. Light incident upon the target produces electron-hole pairs, with the hole migrating to the side of the target, which is addressed by the electron beam. The electron is conducted to ground. The holes accumulate until they are scanned off; the current
348
JAY M. PASACHOFF and DANIEL F. MUZYKA
needed to recharge the diodes is proportional to the incident light intensity. The fact that the spectrum is actually imaged on the vidicon allows the spectrometer to gather data over a wide range of wavelengths. Integration can be achieved internally for 50, 100, 200, 500, or 1000 ms, representing the scan time plus the possible retardations of the scanning electron beam between readings. After the array is scanned, the spectrum is displayed on a Tektronix 7000-series oscilloscope in which a special plug-in module contains half of the spectrometer electronics. The instrument has two gratings. Grating A gives a 4000 A bandpass and is variable in 1000 A steps, while grating B is continuously variable in central wavelength and displays approximately 400 A at 5000 A, and somewhat less with the infrared. A special non-standard grating blazed in the infrared was installed for our eclipse observations. The grating was blazed in the infrared at 10,000 /~ and has about 63~o efficiency at 10,800 A. It was subsequently determined that the actual spectral range displayed during the eclipse was 278/~. Though the spectral sensitivity of the instrument at 10,800 A is only 18~o of the maximum spectral sensitivity in the range of 7000 A to 11,000 A, it is still much better than other methods. The sensitivity corresponds to the convolution of the vidicon quantum efficiency of 7 ~ and the response of the electronics. Its signal at 10,800 A is 3.5 amp A/w. It was of much concern to us that the noise figure of the instrument doubles every 10°C. We had hoped to receive a spectrometer unit that had a liquid/gaseous nitrogen cooling system for the detector, but the technical revisions of the original instrument were not completed by the time of the eclipse and temperature control in Africa was difficult. The importance of the noise figure can more easily be shown by a simple development of the equations involved. The (spot) noise figure F (IRE, 1960) of a system is defined as F _ W.. + W. Wo" ' where W, is the output power attributable to the system and W0, is the output power attributable to the input source. (The gain of a system is defined as W0,/source input power.) It can be shown that this reduces to
F = I + T/ To or
T = (F - 1)To
where T = system noise temperature; To = ambient temperature. The sensitivity or minimum detectable temperature and therefore the minimum detectable flux is given by the rms noise temperature of a system Tsys
Tmi, oC , ~ ,
where T~ysis the system temperature from the previous formula, and t is the integration time. Let us consider an example. If F = 3, To = 290 K, Zy~ by the given formula is 580 K. If one were to integrate for the same length of time and increased the ambient temperature by 10°C to To = 300 K, and therefore double the noise figure to F = 6, the system temperature increases to Zys = 1500 K. This indicates a decrease by a factor of 2.6 in the minimum detectable flux of the instrument. Obviously, this is not a rigorous development of this subject but it does point out our cause for concern at the eclipse. It should be noted that a doubling of a noise figure of approximately 1.5 (To = 290 D), a more realistic figure than F = 3, would yield a decrease in the minimum detectable flux on the order of 4.
Infrared coronal lines DATA
349
COLLECTION
The Williams group was pooled for logistical purposes with other university and research institution groups that were given grants from the National Science Foundation. The National Science Foundation chose to make its prime site in Kenya, where there would be 5 minutes of totality, instead of sites in West Africa, where totality was longer, because of the expected presence of dust in the air in West Africa as a result of the extended drought there. The particular site at the Kenyan oasis of Loiengalani was chosen because it provided a view of the eclipse over water, which would tend to stabilize the images. As it turned out, a dust storm produced about 90~o obscuration on the day of the eclipse at the West African site in Mauritania and ruined many experiments. But the Kenyan site, where we were located, had mostly clear skies. Some small clouds in the sky at totality cleared the sun only 45 seconds before second contact (for a more thorough report of eclipse site activity, see Pasachoff, 1974). Most of our equipment had been shipped to the site at Loiengalani, Kenya; however, we brought the monochromator/vidicon unit of the instrument on the aircraft with us. It was kindly flown from Nairobi to Loiengalani by the Royal Kenya Air Force. The first order of business at the site after arrival was to construct a building to house the spectrometer. As the winds were approximately 60 mph each night, it was useful to construct a solid structure to house the equipment. A floor plan of the building appears as Fig. 4. We also constructed several concrete piers both inside and outside the building to use as stable platforms for the spectrometer and tracking equipment. Unfortunately, because of funding limitations, all the equipment arrived at the last moment before shipping, and was assembled together for the first time on site in Kenya. The effect of the wind-blown sand can be recalled from the example of the air conditioner that we brought with us. Dust got into the contacts of one of the relays and
I 1 3 4 5 6 7 8 9 l0 ii
OSCILLOSCOPE C A L I B R A T I O N LAMPS SPECTROMETER CELESTRON OSCILLOSCOPE V O I C E TAPE R E C O R D E R N O T I O N PICTURE C A M E R A COELOSTAT I N S T R U M E N T A L TAPE RECORDER REFRIGERATOR AIR CONDITIONER
10
WILLIAMS COLLEGE
K-17
FIG. 4. Floor plan of the Williams College spectrometer building at the N.S.F. eclipse site at Loiengalani, Kenya. The equipment in the shed included a film refrigerator (10), two oscilloscopes (1 and 5), the air conditioner, built into the wall (11), the Honeywell instrumental tape recorder (8), the spectrometer unit (3), a box with several calibration lamps (including neon, mercury and argon sources) (7), a Mitchell motion picture camera, used to record the actual pointing of the coelostat (8), a tape recorder for recording operating commands made by the group (6), and the 20 cm Celestron telescope (4). J,V.A. 19/4
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JAY M. PASACHOFF and DANIEL F. MUZYKA
caused it to short out. There is no room here to go into the story of how a slightly smaller 110 V replacement was found and delivered to us in northern Kenya within 24 hours. A major problem at the site was temperature control. Temperatures typically reached 43°C (ll0°F) during the afternoon. Since an internally-cooled spectrometer was not available at the time of the eclipse, we had to try other means of keeping the temperature of the vidicon at the lowest possible level during the eclipse. Besides the air conditioner, we attempted to cool the unit further during the eclipse by placing the remnants of a shipment of dry ice around the unit at the time of the eclipse. (The dry ice had been flown in from Nairobi for the use of other observers.) All of these measures were helpful, but during the actual eclipse a large Lexan window we had installed on the east side of the building had to be removed to eliminate its absorption of incoming infrared radiation. This meant that the building was partially open to the surrounding atmosphere and, despite the cooling efforts, the temperature in the room was approximately 27°C. The temperature in the spectrometer housing itself was probably higher still. This led to a diminished sensitivity for the instrument, as previously discussed. During the eclipse, some of the light from the coelostat was reflected into the Mitchell movie camera. This camera recorded the pointing of the coelostat and therefore the pointing of the telescope. Another diagonal mirror reflected some of the light back to the coelostat mount, to a Nikon camera with cross-hairs, through which the coelostat operator could monitor his adjustments of the pointing. Another camera mounted on top of the Celestron and an adjacent bore-sighted telescope allowed the spectrometer operators inside the house to monitor the tracking as well. Scans of live emission from several spectral lamps (Pen-Ray variety) were taken at intervals before, during and after the eclipse. The lamps were mounted in a box with a slit. Radiation from the box could be reflected into the spectrometer in place of the radiation from the corona by flipping a mirror set between the spectrometer and the Celestron. A filter wheel was placed just before the entrance slit to the spectrometer. Polaroids were mounted in the wheel at 3 polarization angles and one slot was left blank to pass unpolarized light to the spectrometer. The output of the spectrometer unit went to the oscilloscope marked 5 in the floor plan, where the coupling module for the spectrometer was located. The spectra were displayed dynamically on both oscilloscopes. At the eclipse, we recorded the data in FM mode on several channels of a Honeywell 3600C instrumental tape recorder and recorded trigger/timing pulses on other channels. For most of the 5 minutes of totality, we observed in the spectral region that included 10,747 and 10,798. Most of these data were taken of unpolarized radiation. We devoted less than 1 minute at the end of the eclipse to scanning through the spectrum from 7000 h to 11,000 A. Time limitations prevented our observing the 3388 ,~ line of Fe xIn. ANALYSIS
When the data were returned to the U.S., they were first digitized for computer processing. This was carried out at the Tektronix headquarters in Beavertom Oregon. This first digitization referred to as DIG-I, was carried out by playing the analogue tape from the Honeywell instrumental tape recorder into a Tektronix Digital Processing Oscilloscope and recording the data with a PDP 11/25. The intensity level of each scan was sampled at 256 equally spaced points. Each 128 scans were averaged (each scan representing 0.02 s of integration) into "blocks". For the time of the eclipse alone, if the individual scans were left unaveraged, we would have 4 x 10 6 pieces of information. By averaging the data into blocks of 128 scans, the number of points was reduced to approximately 3 x 104. We later discovered that problems arose in the triggering during the first digitization. Because of these problems, some scans and therefore data were lost. Therefore a second digitization of the data was carried out, which we refer to as DIG-2. Each block of
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DIG-2 data represented about 5 s of actual observation time. By February of 1975 most of the digitization problems were corrected and the digital information was transferred from DEC compatible DECTAPE to standard 9 track, 800 bpi magnetic tape and punched cards. We deduced the spectral resolution of the instrument by degrading a section of the Kitt Peak Solar Atlas. We degraded the spectrum between 10,000 A to 10,800 A by moving averages from 0.02 A resolution to resolutions of 1 A to 6 A in order to match the resolution on the vidicon spectra. At the same time, we had available a "solar atlas" made with the eclipse unit. A comparison of the two sets of photospheric spectra showed that the eclipse A-01 unit produced spectra in the region of the two Fe xIn infrared lines of 2 A resolution. Another set of values that had to be determined was actual centers of each spectrum, and what the dispersion of the spectra were. The instrument was centered on 10,800 A, but there is an uncertainty of 10 A in the position of the intensified spot, called the "marker", on the oscilloscope readout. The actual centers of the spectra recorded at the eclipse were determined from calibration source lines. The dispersion corresponded to a total display of 278 A at our infrared wavelength. [re xm] to74rl 3 0 June 1973 F.nllpse at Lolengolnnlo Kenya i
60
h " 1.4 Rn
30
-30
-60 0
j I00
i 200
FIG. 5. An averaged spectrum showing the 10,747 A line of Fe Xlll at a height of 1.4 Ro from the data taken during the 30 June 1973 total solar eclipse.
The spectra were analyzed on the Xerox 530 computing system at Williams College. The upward drifting of successive spectra was corrected by the fitting and subtraction of a first degree polynomial from all of the spectra. Edge effects were dealt with by eliminating the end channels in all arithmetic operations carried out on the spectra. A series of spectra were found to be a good reference baseline, were averaged and then were subtracted from the already "normalized" spectra. An emission feature detected at 10,747 A on the DIG-2 d a t a corresponding to h = 1.4 Ro is shown in Fig. 5. Despite further averaging no feature was detected at 10,798 A and so at this time it is only possible to state a lower limit on the I(10,747 A)/I(10,798 A) ratio of 5/1. Using the Chevalier and Lambert model as adjusted by Finn and Landman for a temperature of T = 2 × 106 K (Fig. 2), this corresponds to an upper limit for the electron density of 3 × 107 for h = 1.4 Ro (Pasachoff and Muzyka, 1975). Flower and Nussbaumer (1974) have recently derived mean values for the electron density corresponding to height of 1.02 Ro from observations of permitted Fe xm lines. The lines in the 200 to 300 • range were observed by Malinovsky and Heroux (1973) from a rocket launched in 1969. Allowing for the difference in heights of observation, the values from the infrared and ultraviolet lines are consistent with each other. The limit for the ratio of intensities we have set gives better agreement between electron densities derived from the Fe xni line ratios and electron densities derived
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JAY M. PASACHOFFand DANIEL F. MUZYKA
from non-Fe xIII methods than did the previous eclipse measurements of Kurt, which Chevalier and Lambert suggested were affected by scattered light. Our results support this conclusion. We are continuing the data reduction by redigitizing the data with a new PDP-11/10 at Williams. We hope that this further data reduction will sufficiently improve the signalto-noise ratio to detect the 10,798 line at this and other coronal heights, and thus give us more values of the electron density.
ACKNOWLEDGEMENTS We thank Mr. J. Philip Schierer and Tektronix, Inc. for lending us the spectrometer and for giving various other types of assistance in several stages of this project. Dr. William Weir of Reed College and Tektronix has kindly assisted with programming for setting up the Digital Processing Oscilloscope at Tektronix. Messrs. Stuart N. Vogel '75 and Dan Stinebring '75 also participated in the expedition, and were a vital part of all phases of this experiment in Africa. Mr. Mark Chaffee '77 made the measurements of the pointing. We also thank many other individuals at Williams College for their help. Prof. Stuart J. B. Crampton was of valuable aid in transferring digital data for us. Prof. Howard Stabler designed an FM modulator/ demodulator, which was built by Mr. Emile Ouellette. Messrs. Joseph McCann and Albert St. Pierre expertly constructed mountings, packing crates, and other necessities. We thank Professor Donald H. Menzel of the Center for Astrophysics for useful conversations and for introducing us to the techniques of eclipse observing that he has perfected over the years. Donald A. Cooke '75 assisted with the equipment for this experiment in Williamstown, and then went on to a site in West Africa with Professor Menzel to carry out another spectrographic experiment. We are grateful to him for taking on this responsibility. We are especially thankful to Mr. Ronald R. La Count, U.S. Solar Eclipse Coordinator at the National Science Foundation, Dr. G. William Curtis, Director of the N.S.F. expedition, and Mr. Eugene Prantner and the entire support staff of the Expedition from the National Center for Atmospheric Research in Boulder. We also thank Dr. Harold Lane of the N.S.F. for his encouragement of this project. We are grateful to Mr. Bud Corbin of the Honeywell Corporation for making their 5600C instrumental tape recorder available to us both in Africa and in Oregon. We also acknowledge with thanks the assistance of Mark Armstead, Inc. of Hollywood, California, for lending us the Mitchell movie cameras used for some of the time-lapse movies, and of Messrs. Dennis Kane and Sidney Platt of the National Geographic Society for lending us an Arriflex movie camera for other time-lapse work. We wish to thank Mrs. Mary G. Smith, and others at the National Geographic Society for their encouragement over and above the financial support that we have acknowledged below. We also appreciate the assistance of the Photographic Division of National Geographic, especially Messrs. Robert Gilka and John Fletcher and the people in the black and white and color labs who worked very hard on technically tricky negatives, for providing camera equipment and for handling all of the film processing and printing. We thank Dr. A. Keith Pierce and Mr. Larry Testerman of the Kitt Peak National Observatory for providing us with their solar Atlas in advance of publication. Dr. Robert F. Howard of the Hale Observatories and Dr. Stephen A. Schoolman, then of Caltech and now of the Lockheed Solar Observatory, were of great assistance in the definition stage of this experiment at the time the spectrometer first became available to us. We would like to thank the Computer Center at Williams for providing D.F.M. with time off from his position and for providing large amounts of time on the computer. We were supported in part by National Science Foundation grants GP37187 and MPS 74-23531, and grants from the National Geographic Society and the Research Corporation. We would also like to acknowledge support from a grant by the Alfred P. Sloan Foundation to Williams College, and from the President's Discretionary Fund at Williams. The Government of Kenya deserves many thanks for the hospitality and services it provided for us and for the entire expedition.
REFERENCES Bahcall, J. and Wolf, R. A. (1968) Ann. Astrophys. 29, 343. Byard, P. and Kissell, K. (1971) Solar Phys. 21, 351. Chevalier, R. and Lambert, D. (1969) Solar Phys. 10, 115. Dumont, J.-P. and Perche, J.-C. (1964) 12me Coll. Int. Astrophys. Libge p. 86. Eddy, J. A., Firor, J. W. and Lee, R. H. Bull. Am. Astron. Soc. 72. Eddy, J. A., Lee, R. and Emerson, J. (1973) Solar Phys. 30, 351. Eddy, J. A. and Malville J. (1967) Astrophys. J. 150, 289. Edl6n, B. (1942) Z. Astrophys. 22, 30. Finn, G. D. and Landman, D. (1973) Solar Phys. 30, 381. Firor, J. and Zirin, H. (1962) Astrophys. J. 135, 122. Fisher, R. and Pope, T. (1971) Solar Phys. 20, 389. Flower, D. R. and Nussbaumer, H. (1974) Astron. Astrophys. 31, 353. IRE. Standards on Receivers, Definition of Terms, Proc. IRE 40, 1681. Kurt, V. G. (1962) Soviet Astron. 6, 349.
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Kurt, V. G. (1963) Soviet Astron. 6, 620. Lyot, B. (1939) M.N.R.A.S. 99, 580. Malinovsky, M. and Heroux, L. (1973) Astrophys. J. 181, 1009. Malville, J. (1967) Astrophys. J. 148, 299. Pasachoff, J, M. (1974) unpublished report. Pasachoff, J. M. and Muzyka, D. F. (1975) Bull. Am. Astron. Soc. 7, 409, Ratier, G. and Rozelot, J.-P. (1971) Solar Phys. 23, 394. Shklovskii, J. S. (1965) Physics of the Solar Corona, Pergamon Press, Oxford. Van Regemorter, H. (1962) Astrophys. J. 136, 906. Wl6rick, G., Dumont, J.-P. and Perche, J.-C. (1963) in The Solar Corona, Ed,, J. W. Evans, p. 177, Academic Press, New York. Zirin, H. (1966) The Solar Atmosphere, Blaisdell. Zirin, H. (1970) Solar Phys. 11, 497.