ELECTRON TEMPERATURE AND SPEED ... - IOPscience

The Astrophysical Journal, 599:596–603, 2003 December 10 # 2003. The American Astronomical Society. All rights reserved. Printed in U.S.A.

ELECTRON TEMPERATURE AND SPEED MEASUREMENTS IN THE LOW SOLAR CORONA: RESULTS FROM THE 2001 JUNE ECLIPSE Nelson L. Reginald Catholic University of America at NASA Goddard Space Flight Center, MC 682, Greenbelt, MD 20771; [email protected]

O. C. St. Cyr and Joseph M. Davila NASA Goddard Space Flight Center, MC 682, Greenbelt, MD 20771

and Jeffrey W. Brosius Catholic University of America at NASA Goddard Space Flight Center, MC 682, Greenbelt, MD 20771 Received 2003 July 25; accepted 2003 August 18

ABSTRACT We present measurements of electron temperature and bulk flow speed low in the solar corona derived from white-light spectra obtained during the total solar eclipse of 2001 June 21. Observations were obtained at two locations in the solar corona, one within a helmet streamer at the east limb and the second in a streamer cluster in the southwest. Both points were at an altitude of about 1.1 R
from the solar center. The east limb and southwest locations yielded electron temperatures of 0:96  0:05 and 1:2  0:2 MK and bulk þ443:0 1 flow speeds of 72:0þ281:0 72:0 and 257:0257:0 km s , respectively. These measurements are unique in that they simultaneously provide both the electron temperature and its bulk flow speed; few previous measurements of electron parameters in the corona are available. The results presented here demonstrate the potential for this technique: if the instrument were used with a coronagraph, it would provide routine synoptic maps of electron temperature and bulk flow speed. Subject headings: eclipses — Sun: atmospheric motions — Sun: corona — Sun: UV radiation both the temperature and the bulk flow speed of electrons, simultaneously from multiple locations, low in the Sun’s corona using a ground-based instrument in conjunction with a total solar eclipse. Knowledge of the density, temperature, and speed parameters for various constituents of the solar wind provides essential constraints to solar wind models. It is important to obtain measurements for the electrons, since they may be accelerated and/or heated at different rates or by different mechanisms than the heavier ions. The Sun’s white-light corona is visible only during a total eclipse or by means of coronagraphic techniques. Its K-coronal component, formed by Thomson scattering of photospheric radiation by free electrons, dominates the F component (scattering by dust) and E component (line emission) at low altitudes. Thermal broadening due to the high electron temperature in the corona obliterates almost all of the Fraunhofer absorption lines and most other narrow features in the Sun’s visible-light continuum spectrum. Measurements made by Grotrian (1931) during a total eclipse showed some hint of weak residual depressions in the Thomson-scattered coronal spectrum at the wavelengths of the strongest absorption lines. The basic description of coronal brightness due to Thomson scattering is well known (e.g., van de Hulst 1950; Billings 1966; Hundhausen 1993) and will not be repeated here. Cram (1976) developed a detailed theoretical description of the Thomson scattering process in the solar corona. In Cram’s theory, the intensity of the scattered light is calculated as an integral over a stationary Maxwellian distribution function that is temperature dependent, which in turn determines the shape of the K-coronal spectrum. Cram suggested that the temperature of the coronal electron could be determined to within 0.2 MK if the ratio of the intensities ˚ could be measured to an accuracy of at 3900.0 and 4100.0 A

1. INTRODUCTION

Coronal temperatures have been determined by numerous authors based on line emission from ions (e.g., Habbal, Esser, & Arndt 1993; Guhathakurta, Fisher, & Altrock 1993; Brosius et al. 1997; Mason et al. 1997; Kohl et al. 1995) and semiempirical modeling (e.g., Guhathakurta et al. 1999). Current knowledge of the solar wind speed close to the Sun comes from several different techniques also based on observing emission from ions: Doppler dimming of UV spectral features (e.g., Strachan et al. 1993) and visible imaging spectroscopic coronagraphy (e.g., Brueckner et al. 1995). Interplanetary scintillation radio measurements (e.g., Grall et al. 1996) also provide a measure of the solar wind speed using electrons. The only other measurement of electron parameters low in the solar corona is that described by Fineschi et al. (1998) using Ly
profiles from observations with the Ultraviolet Coronagraph Spectrometer (UVCS) aboard the Solar and Heliospheric Observatory (SOHO) satellite. Much of the interesting physics relevant to accelerating the solar wind and initiating coronal mass ejections occurs in the lowermost region of the corona, from 1.1 to 2.0 R
. The Extreme Ultraviolet Imaging Telescope (EIT) on SOHO (Delaboudiniere et al. 1995) provides images in four narrow wave bands, each corresponding to different formation temperature (e.g., Moses et al. 1997), with a field of view extending to a few tenths of a solar radius above the limb. The UVCS observations begin at 1.5 R
(Kohl et al. 1997), while observations with the Large Angle and Spectrometric Coronagraph (LASCO) cover heights up to several tens of solar radii (e.g., Socker et al. 1998). We report here the results of a novel spectroscopic technique described by Reginald & Davila (2000) to measure 596

CORONAL ELECTRON TEMPERATURE AND SPEED 1.0%. These wavelengths correspond to peaks and dips (e.g., Fig. 3 in Cram 1976) in the K-coronal spectrum because of the residual absorption features of the calcium H and K lines and the G band (a collection of absorption features predominantly due to iron). Cram showed that the shape of the visible light continuum could be used to discriminate between various temperatures of a static, spherically symmetric, isothermal model corona. Ichimoto et al. (1996) used Cram’s technique to measure the electron temperature during the total solar eclipse of 1994 November. They used a slit spectrograph with a cooled CCD to obtain spectra at two altitudes in a coronal hole and in a coronal streamer. Clouds caused extinction and produced additional scattered light in their observations, but they were able to derive electron temperatures of 1.5– 1.7 MK for the streamer and 0.9–1.1 MK for the south polar coronal hole. Reginald & Davila (2000) extended Cram’s model of the K-coronal spectrum to include effects of both the temperature and bulk flow speed of the coronal electrons due to the solar wind and designed an experiment to measure these two quantities simultaneously from multiple locations in the solar corona. In x 2, we provide a summary of the theory. In x 3, we describe our instrument and observations, in x 4 we present our results, in x 5 we discuss implications of this investigation, and in x 6 we summarize our conclusions.

2. THEORY

The solar photospheric spectrum at visible wavelengths is replete with absorption lines (Fig. 1). Among the deep and conspicuous absorption lines, which were labeled by Fraunhofer, are the calcium H and K lines at 3968.5 and ˚ , respectively. These lines are very important in 3933.7 A what follows. The K corona is a result of Thomson scattering of photospheric radiation by coronal electrons. Unlike the photospheric spectrum, the K-coronal spectrum is smooth and continuous because of thermal Doppler broad˚ will be ening. For example, incident radiation of 4000.0 A ˚ thermally Doppler broadened by 70.0 A in a 1.0 MK

Fig. 1.—Extraterrestrial solar irradiance spectrum obtained with the ground-based Fourier transform spectrometer at the McMath/Pierce Solar Telescope at Kitt Peak, Arizona, corrected for wavelength-dependent absorption in the Earth’s atmosphere.

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Fig. 2.—Modeled K-coronal intensity spectra as a function of wavelength for isothermal coronal temperatures of 0.5, 1.0, 1.5, and 2.0 MK and an electron bulk flow speed of 0.0 km s1 for lines of sight at 1.1 and 1.3 R
. ˚. The temperature-insensitive nodes are at 3987.0, 4234.0, and 4411.0 A

plasma. This smooths the narrow absorption lines of the incident photospheric radiation into a continuum with weak depressions in the vicinity of the deep Fraunhofer lines and served as one of the first clues for a hot corona. Figure 2 shows the modeled shape of the K-coronal spectrum for photospheric light scattered off the coronal electrons situated along lines of sight at 1.1 and 1.3 R
from the solar center for a series of isothermal coronal temperatures of 0.5, 1.0, 1.5, and 2.0 MK. It is evident from this modeled spectrum that the narrow absorption features of the incident photospheric radiation (Fig. 1) have been heavily smoothed by thermal Doppler broadening. The scattered intensity spectrum is smooth, with shallow depressions that roughly correspond to absorption features in the incident spectrum. This smoothing increases with increasing temperature and gives rise to temperature sensitive antinodes (peaks and dips) indicated in Figure 2. The existence of temperature-insensitive nodes (crossing points) in the Kcoronal intensity spectra, along with the wavelength independence of these nodes for lines of sight at various heights above the solar limb, was first reported by Cram (1976). We can determine the coronal electron temperature at a given height by fitting the measured K-coronal spectrum with modeled spectra at different temperatures, using the nodes as anchors. The hot solar corona is not static but continuously flows out into interplanetary space. This is known as the solar wind. Assuming that the coronal electrons travel away from the Sun at the solar wind speed, the photospheric radiation incident on these electrons is redshifted in the frame of reference of the electrons. Furthermore, for Thomson scattering, the photospheric radiation scattered off these electrons is wavelength independent in the reference frame of the electrons, and they scatter a fraction of this incident radiation along the line of sight to the observer. An additional interesting property of the modeled K-coronal spectra was revealed by including a constant radial bulk outflow speed for the coronal electrons (Reginald 2000). This causes the temperature insensitive nodes in Figure 2 to shift in wavelength positions with increasing electron bulk flow speed but maintains the height independence of the wavelength

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3. INSTRUMENT AND OBSERVATIONS

Fig. 3.—Modeled K-coronal spectral intensity as a function of wavelength for isothermal coronal temperatures of 0.5, 1.0, 1.5, and 2.0 MK and an electron bulk flow speed of 400.0 km s1 for a line of sight at 1.1 R
. The ˚. temperature insensitive nodes are at 3990.0, 4239.0, and 4416.0 A

positions of the nodes. Figures 3 and 4 show model Kcoronal intensities for different isothermal coronal temperatures for a line of sight at 1.1 R
from the solar center, with assumed radial electron bulk flow speeds of 400.0 and 800.0 km s1. Comparing Figures 3 and 4 with Figure 2, which is the model K-coronal intensity for zero bulk flow speed, it is evident that the temperature insensitive nodes increase in wavelength positions with increasing electron bulk flow speeds. For example, the magnitude of this redshift on the ˚ K-coronal spectrum for an incident radiation of 4000.0 A on coronal electrons with an outflow speed of 300.0 km s1 ˚ . The existence of these temin a 1.0 MK corona is 4.0 A perature insensitive nodes, combined with their wavelength shifts due to the bulk flow speed of the electrons, enables us to measure both the temperature and the bulk flow speed of these electrons from measurements of K-coronal spectra (Reginald 2000).

Fig. 4.—Modeled K-coronal spectral intensity as a function of wavelength for isothermal coronal temperatures of 0.5, 1.0, 1.5, and 2.0 MK and an electron bulk flow speed of 800.0 km s1 for a line of sight at 1.1 R
. The ˚. temperature-insensitive nodes are at 3994.0, 4244.0, and 4421.0 A

The Multi-Aperture Coronal Spectrometer (MACS) was designed to obtain simultaneous spectra from multiple locations in the corona during a total solar eclipse. A commercial 12.0 inch, f/10 Schmidt-Cassegrain telescope produced the images. The telescope f number was reduced to f/6.3 using a focal reducer. The plate scale of the telescope with the focal reducer in place was 0>107 lm1, so that each of the 200.0 lm diameter fibers located at the focal plane sampled 21>4 of the solar corona. The 6:0  6:0 cm2 transmission grating had a ruling density of 600.0 lines mm1. All the optical components had antireflective coating to ˚ . The camera enhance the transmission below 4000.0 A incorporated a thermoelectric cooler capable of cooling to 50 C below ambient, and a back-thinned SITE CCD of size 512  512 pixel2, with each pixel of dimension 24  24 lm2. The readout time was 10.0 s. The spectral and spatial reso˚ pixel1 and 5>0 pixel1, respectively, lutions were 2.46 A ˚ at the CCD with a spectral range from 3500.0 to 4600.0 A detector. The details of the instrument design can be found in Reginald & Davila (2000). Although MACS has 25 locations for fibers in the focal plane, only six were used at the 2001 June eclipse. The MACS instrument was deployed to observe the 2001 June 21 total solar eclipse at a location near Lusaka, Zambia. A description of the eclipse can be found in Espenak & Anderson (1999). We obtained five exposures, with durations of 3.75, 11.25, 15.0, 33.75, and 101.25 s. All five exposures were added to increase the signal-to-noise ratio in the present analysis. Figure 5 shows the low corona in white light a few hours after the eclipse as observed in polarized brightness by the Mk 4 K coronameter at Mauna Loa Solar Observatory (Fisher et al. 1981). The pattern of the fiber apertures in the MACS focal plane has been overlaid on the Mk 4 image. Since the moon’s shadow drifted across fibers 16 and 20 in Figure 5, only three fibers (4, 8, and 12) located at 1.1 R
from the solar center provided useful spectra. Based on Mk 4 and LASCO observations, fiber 12 was located near the equatorial edge of an east limb helmet streamer and fiber 4 in the southwest in a cluster of small streamers. The observed data required wavelength calibration to determine the dispersion relationship and intensity calibration to account for modulation due to atmosphere and instrumental optics. The wavelength calibration of the instrument was derived from a prominence spectrum obtained in fiber 20 during the 15.0 s exposure before it slipped into the lunar shadow region. This spectrum provided clearly identifiable emission lines that were used to determine the wavelength scale. The average linear disper˚ matches very well with the design sion of 2:40  0:04 A expectation based on the optical parameters for the spectrograph. The intensity calibration required a more complicated procedure. First a ‘‘ terrestrial ’’ (observed) photospheric spectrum ITP ðÞ was obtained by pointing the MACS instrument at the Sun through a pinhole with a known aperture. The pinhole was necessary to limit the solar flux entering the instrument and avoid damage to the optics. The corresponding ‘‘ extraterrestrial ’’ (true) photospheric spectrum IEP ðÞ (Fig. 1) was obtained at the McMath/Pierce Solar Telescope at Kitt Peak, Arizona, through a Fourier transform spectrometer, corrected for wavelength-dependent

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Fig. 5.—Locations of the six MACS fibers during the solar eclipse of 2001 June 21. Fibers 16 and 20 drifted into the lunar shadow during observation, so that only fiber 4, 8, and 12 provided coronal observation. The coronal image was taken at 17:33 UT while the observation through MACS was taken between 13:09 and 13:12 UT. Coronal image courtesy Mauna Loa Solar Observatory and the High Altitude Observatory.

absorption in the Earth’s atmosphere (Kurucz et al. 1984). This spectrum is considered to be a reliable, well-calibrated standard. The resolution of this extraterrestrial photospheric spectrum was reduced to match the MACS resolution. The ratio between the terrestrial photospheric spectrum ITP ðÞ and extraterrestrial photospheric spectrum IEP ðÞ was then used to derive the intensity calibration curve. This was used to transform the terrestrial (observed) coronal spectrum ITC ðÞ to the extraterrestrial (true) coronal spectrum IEC ðÞ as follows: IEC ðÞ ¼

½ITP ðÞ

IEP ðÞ=IEP ð4000Þ I C ðÞ : ð1Þ  CT P 2 2   =½IT ð4000Þ  4000  IT ð4000Þ

The first term on the right-hand side of equation (1) is the calibration curve. For convenience, all measured intensities ˚ . Diffraction effects due to have been normalized at 4000.0 A the pinhole introduced wavelength dependence in the size of the solar image; we accounted for this by the 2 factor in equation (1). Note that the pinhole was used only during the

terrestrial photospheric measurements ITP ðÞ and not during the terrestrial coronal measurements ITC ðÞ. Figure 6 shows a terrestrial coronal spectrum ITC ðÞ and its derived extraterrestrial coronal spectrum IEC ðÞ. In all of the images that were used for calibration and data analysis, the corresponding dark images were subtracted. These dark images were taken soon after totality, with the same exposure times as those of the eclipse images. The flat-fielding of the individual images was performed by exposing the camera to the sky through a diffuser and assuming the sky to be of uniform brightness.

4. RESULTS

We identified the K-coronal model with the electron bulk flow speed and temperature that best fits with the extraterrestrial coronal data IEC ðÞ. We inverted the K-coronal models of Reginald & Davila (2000) for a range of electron temperatures (from 0.5 to 1.0 MK in increments of 0.05

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Fig. 6.—Terrestrial (observed) coronal spectrum ITC ðÞ and its derived extraterrestrial (true) coronal spectrum IEC ðÞ, for fiber 12. This transformation was made using the calibration curve shown by the first term on the right-hand side of eq. (1).

MK, and from 1.0 to 2.0 MK in increments of 0.1 MK) at each of the eight values for bulk flow speed (from 0.0 to 700.0, in increments of 100.0 km s1) to obtain a relationship between modeled K-coronal intensity and modeled temperature for each value of the wavelength (from 3500.0 ˚ at 1.0 A ˚ resolution) and each value of the bulk to 4500.0 A flow speed. A polynomial fit of modeled K-coronal intensity as a function of modeled temperature was derived at each of the different combinations of wavelength and bulk flow speed. We restricted the analysis to the wavelength ranges that corresponded to the polynomials that reproduced the modeled K-coronal intensity for all of the above model temperatures to an accuracy of 1.0%. An additional constraint was that the polynomials provided a unique solution. See Figure 2, where we show model K-coronal spectra calculated at four different temperatures with zero bulk flow speed. The curves converge at the nodes located at 3987.0 ˚ , where they cross over each other and, and 4234.0 A hence, produce multiple temperatures for a given modeled K-coronal intensity value. We refer to the wavelength ranges that satisfy the above criteria as range 1. The observed spectra revealed wavelength ranges in which the noise-to-signal ratio was less than 1.0%: we refer to this as range 2. Cram (1976) showed that such a noise level would yield an uncertainty of 0.2 MK on temperature measurements, and Reginald & Davila (2000) demonstrated that it would translate to an uncertainty of 200.0 km s1 in bulk flow speed measurements. In what follows, we focus on wavelength ranges corresponding to the overlap of ranges 1 and 2. For the spectrum observed through fiber 12, these wavelength ranges include 3910.0–3969.0 and ˚ . Unfortunately, for fiber 4, we obtained no 4268.0–4356.0 A better than 2.0% noise-to-signal ratio, so range 2 for this fiber incorporated this constraint. As a consequence, results derived for fiber 4 yielded larger error bars. Figures 7 and 8 show results for the above two wavelength ranges for the spectrum observed through fiber 12. The plus (+) marks in these figures indicate the extraterrestrial coronal spectrum IEC ðÞsuperposed on our K-coronal models for electron bulk flow speed of 100.0 km s1. The corresponding temperatures derived for each of the

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Fig. 7.—Extraterrestrial coronal spectrum IEC ðÞ from fiber 12, (see Fig. 6 for shape and Fig. 5 for location), shown as pluses (+) across the wave˚ . Solid lines: Modeled K-coronal intensities, length range 3910.0–3969.0 A from bottom to top, at electron temperatures of 0.5–1.0 MK in steps of 0.05 MK and 1.0–2.0 MK in steps of 0.1 MK, at a wind speed of 100.0 km s1. ˚ . The crosses () show the These have been normalized to unity at 4000.0 A temperatures determined at observed wavelength positions within this wavelength range.

observed wavelength positions are indicated by crosses (). Similarly, temperatures were derived for electron bulk flow speeds between 0.0 and 700.0 km s1. Results for both wavelength ranges are displayed in Figure 9. The ‘‘ best ’’ temperature and flow speed occurs at the point where the two lines intersect. For fiber 12, this occurs at 0:96  0:05 MK and 1 72:0þ281:0 72:0 km s . Associated measurement uncertainties were derived from random and statistical noise as follows. The random noise level was determined to be 16.0 counts (65,536 counts maximum). This number was derived in the following manner using three dark images taken after totality. At each pixel location along the spectral direction, the dark counts were first averaged over the spatial extent of the three individual dark images. Then these counts were averaged at each pixel position along the spectral extent, and the corresponding standard deviations were calculated. The random noise level was taken to be twice the average of these standard deviations, thus accounting for the noise in

˚ Fig. 8.—Same as Fig. 7, but in the wavelength range 4268.0–4356.0 A

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Fig. 9.—Temperatures determined in the two wavelength ranges for a sequence of model bulk flow speeds. The average temperature is shown for each of the eight bulk flow speeds from 0.0 to 700.0 km s1 in intervals of 100.0 km s1. Dashed and solid lines: Least-squares fits through those points. Dotted lines: Range of measurement uncertainties. The intersection of the lines represents the temperature and bulk flow speed where data from the two wavelength ranges agree. This plot yields a coronal electron temperature of 0:96  0:05 MK along with electron bulk flow speed of 72:0þ281:0 72:0 km s1 for fiber 12. The shaded region encompasses the upper and lower limits for temperature and bulk flow speed, with the constraint that the flow speed remains positive.

both the exposed and dark images used in this analysis. In the absence of a calibration source, the photon statistical error was determined using the parameters associated with the camera. Here the square of the Poisson error reflecting the statistical error was determined by calculating the number of electrons (300,000 full well depth) associated with the net (observed counts minus the dark counts) measured digital counts (65,536 maximum) at the observed wavelength positions. The intensity calibration process described in x 3 effectively removes systematic measurement errors from our analysis. For completeness, we show in Figure 10 a plot of the extraterrestrial coronal spectrum IEC ðÞ of fiber 12 as well as the modeled K-coronal spectrum for isothermal temperature of 0.96 MK and a bulk flow speed of 72.0 km s1. Here ˚ . The the data have been normalized to unity at 4000.0 A agreement between the two is quite good. The two ranges bounded by the vertical lines (marked in gray) correspond to the two ranges described in Figures 7 and 8, where models were matched with the observed data. By the same procedure described above, the electron temperature and bulk flow speed determined for fiber 4 were 1 1:2  0:2 MK and 257:0þ443:0 257:0 km s , respectively. The uncertainties are greater here because the noise level was nearly twice that of fiber 12. Figure 11 shows the extraterrestrial coronal spectrum IEC ðÞ of fiber 4 as well as the modeled K-coronal spectrum for isothermal temperature of 1.2 MK and a bulk flow speed of 257.0 km s1. For fiber 8, the above procedure failed to converge, although the original spectrum did qualitatively match the model.

5. DISCUSSION

The technique described above allows one to derive the electron temperature and its bulk flow speed using the

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Fig. 10.—Extraterrestrial coronal spectrum IEC ðÞ derived for fiber 12 (see Fig. 6 for shape and Fig. 5 for location) plotted along with the modeled K-coronal spectrum for isothermal coronal temperature of 0.96 MK and electron bulk flow speed of 72.0 km s1 (see Fig. 9). Here the plots have been ˚ . The two shaded ranges correspond to normalized to unity at 4000.0 A the two ranges described in Figs. 7 and 8, where models were matched with the observed data (overlap between ranges 1 and 2).

white-light measurements of the K-coronal spectrum at different locations in the solar corona. The measured coronal spectrum actually consists of three components: the ‘‘ true ’’ electron-scattered K corona, the F corona produced by scattering off interplanetary dust between the Sun and the observer, and the E corona produced by emission lines. There is no direct evidence of emission lines in the spectra from fibers 4, 8, and 12. At the altitudes of the fiber locations, the E corona is typically 2 orders of magnitude fainter than the K corona, and the F corona 1 order of magnitude fainter than the K corona (Zirin 1988); both of these have

Fig. 11.—Extraterrestrial coronal spectrum IEC ðÞ derived for fiber 4 (see Fig. 5 for location) along with the modeled K-coronal spectrum for isothermal coronal electron temperature of 1.2 MK and electron bulk flow speed of 257.0 km s1. Also plotted are the models for temperatures of 1.0 and 1.4 MK, which are the upper and lower bounds for the temperature. It is apparent that the models for these upper and lower bounds envelope the observed data in the two wavelength ranges where comparisons were made between the observed results and the K-coronal models. The two shaded ranges show these two wavelength ranges. Here the plots have been ˚. normalized to unity at 4000.0 A

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been neglected in our analysis. For coronagraphic observations, contributions from the F corona can be easily removed through the use of polarized brightness measurements through three different orientations of the polarization axis. We did not have sufficient time to perform such measurements during the eclipse (216.0 s). A direct measurement of the coronal electron temperature was derived by Fineschi et al. (1998) from K-corona spectral line profiles of H i Ly
measured with the UVCS instrument on board SOHO, accounting for the stray light profile of the measured H i Ly
profiles from experiments generated in the laboratory. The derived temperature was 1:1  0:3 MK in a streamer at 2.7 R
. Based on line ratio techniques involving different ionization states of Fe, Parenti et al. (2000) used SOHO/CDS and SOHO/UVCS spectral-line data to measure the coronal temperature at heights of 1.02–1.19, 1.58, and 1.60 R
. The temperature measured at midlatitude at 1.58 and 1.60 R
was 1:28  0:02 MK, while the temperatures measured along an equatorial streamer and equator (at the same two above heights) were 1:24  0:01 and 1:23  0:01 MK, respectively. For heights from 1.02 to 1.19 R
, the temperatures measured ranged from 1.08 to 1.25 MK for an equatorial streamer and 1.23 to 1.50 MK for a midlatitude streamer. Wilhelm, Inhester, & Newmark (2002) also applied line ratio techniques to SOHO (SUMER, LACSO/C1, EIT) data to measure electron temperatures of 0.9 MK very close to the limb at the equator and  1.5 MK at  1.2 R
at 35 latitude. Ichimoto et al. (1996) applied Cram’s (1976) theoretical prescription to eclipse spectra obtained 1994 November 3 in Putre, Chile, to measure temperatures of 1.71 MK at 1.1 R
in an east limb coronal streamer and 1.07 MK at the base of a south pole coronal hole. Marsch et al. (1999) determined proton temperatures of 0.1 to 0.2 MK at the low solar polar corona at less than 1.06 R
using SOHO/SUMER spectra. These measurements of electron temperature do not seem to show a pattern with increasing coronal heights: most authors report values between 1.0 and 2.0 MK within a radius of 2.0 R
and temperatures less than 1.0 MK close to the limb. Our measurements of 0.96 and 1.2 MK at 1.1 R
close to the equator are consistent with the above values. Strachan et al. (1993) used Doppler dimming analysis of resonantly scattered H i Ly
obtained from an ultraviolet light coronagraph flown on a sounding rocket to derive pro1 at 2.0 R with a confiton velocities of 217:0þ34:0
64:0 km s dence level of 67%. Using a two-fluid model constrained by density profiles inferred from space-based SPARTAN 20101 and ground-based Mauna Loa white light coronagraph observations together with in situ interplanetary measurements, Habbal et al. (1995) measured a solar wind speed of 4.0 km s1 and electron temperature of 0.8 MK at 1.2 R
for both an ambient coronal hole and a dense structure. At 1.5 R
, these values increased to 20.0–100.0 km s1 and 1.0 MK, respectively. Interplanetary scintillation techniques provide reliable measurements of electron bulk flow speeds at higher altitudes in the corona. Grall et al. (1996) measured an electron bulk flow speed of 600.0 km s1 close to 6.0 R
. Using the Charge, Element, and Isotope Analysis System/Mass Time-of-Flight (CELIAS/MTOF) proton monitor instrument on SOHO, Ipavich et al. (1998) measured the proton bulk flow speed at 220.0 R
of 421:0  76:0 km s1. Based on SOHO-Ulysses quadrature observations, Poletto et al. (2002) measured low-latitude

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proton bulk flow speeds of 100.0 and 150.0 km s1 at 3.5 and 4.5 R
, respectively, in coronal holes and 35.0 and 60.0 km s1 at the same respective heights along streamers. Taking the proton bulk flow speed as a measure of the solar wind speed indicates that the solar wind speed increases from 100.0 km s1 at 1.5 R
to 600.0 km s1 at 6.0 R
, consistent with expectations of an accelerating wind in that height range. In this context, our measurements of electron bulk flow speeds of 72.0 and 257.0 km s1 at 1.1 R
might be greater than expected. However, this underscores the importance of independently measuring the bulk flow speeds of electrons, protons, and ions: only by doing so can we obtain valuable information regarding acceleration mechanisms in the source region of the solar wind. MACS could be used to determine the acceleration of the solar wind by measuring the electron flow speed in multiple fibers at different heights in the corona. Ichimoto et al. (1996), who also used Cram’s (1976) prescribed technique, estimated an electron bulk flow speed increase 80 km s1 between 1.1 and 2.0 R
along a streamer based in slit spectra obtained during the total solar eclipse of 1994. This measurement was based on the ratio of the observed K-coronal ˚ , which are insensitive to intensities at 3990.0 and 4218.0 A the electron temperature but sensitive to the Doppler shifts. The approximate relation shown in their paper is based on the assumption that the slope of this ratio for different bulk flow velocities is the same at all heights in the corona. However, Reginald & Davila (2000) have shown that the slope of this intensity ratio against bulk flow speed of the coronal electrons is different for different coronal heights (see their Fig. 4). There are several potential improvements to the technique described in this paper. First, it is possible that a more realistic, multithermal coronal model might yield better fits to the observations. The K-coronal spectrum along a line of sight is heavily dependent on the electron density distribution close to the intersection of the line of sight with the plane of the solar limb because of the rapid decrease in electron density with radius. However, as discussed in Reginald (2000), a multithermal corona would be more plausible if there are streamers crossing the line of sight contributing to enhanced coronal electron density at those locations. Second, if we could enhance the signal-to-noise ratio, we would reduce the uncertainty in the bulk flow speed measurements. This would be achieved, for example, by incorporating a reflection grating for wavelength dispersion instead of a transmission grating. The MACS instrumental concept could be incorporated behind a space-borne coronagraph mounted on a satellite or a ground-based coronagraph located at a high altitude. This would enable much for frequent observations and allow sufficient time to obtain the polarization measurements needed to remove the F-corona contributions. This method represents a novel approach to simultaneously and globally measure both the electron temperature and its bulk flow speed in the Sun’s corona. 6. SUMMARY

We designed, constructed, and deployed the Multiaperture Coronal Spectrometer (MACS) to measure white-light K-coronal intensities simultaneously at different locations in the corona during the total solar eclipse of 2001 June 21. The data were calibrated and compared with modeled

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K-coronal intensities in order to simultaneously derive the electron temperature and its bulk flow speed. We obtained an electron temperature of 0:96  0:05 MK and a bulk flow 1 near the edge of a helmet streamer speed 72:0þ281:0 72:0 km s and an electron temperature of 1:2  0:2 MK and a bulk 1 in a cluster of small streamers. flow speed 257:0þ443:0 257:0 km s Both locations were at heights 1.1 R
from the solar center. This powerful technique ultimately can be applied to

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measure the electron acceleration in the source region of the solar wind through the use of multiple fibers at different heights. M. Guhathakurta and A. Harper provided valuable assistance obtaining these data. N. L. R., O. C. S., and J. W. B. acknowledge partial support from NASA grant NAG59502.

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