The identification of Feix and Nixi in the solar corona

The levels of the configuration 3s ²3 _p ⁵3d of Fe ix have been experimentally determined from their combinations with 3s3 _p ⁶3d ³ D in the region 300–400 Å. Wavelengths can now be accurately predicted for all transitions within 3s ²3 _p 53^d, and eight of these can be identified with coronal lines from 2042 to 4585 Å. Also, identifications of solar lines from 171 to 245 Å with electric-dipole and magnetic-quadrupole transitions to the ground state, 3s ²3p ^{6 1} S, are confirmed and extended. Solar identifications with corresponding transitions in Ni xi, both within 3s ²3 _p ⁵3d and to the ground state, are proposed on the basis of a short extrapolation.     

The forbidden transitions within 3s 2 3p 5 3d of Fe ix and Ni xi and 3s 2 3p 4 3d of Fe x and Ni xii

Improved level values for 3s ² 3p ⁵ 3d and 3s ² 3p ⁴ 3d of Fe ix and Fe x are derived from the observed forbidden lines with special regard to recent accurate measurements in the vacuum-ultraviolet part of the coronal spectrum. A procedure for estimating expected relative intensities is proposed and used for an additional check on the consistency of the identifications, which now comprise 12 lines of Fe IX and 11 lines of Fe x. Finally, by use of a suitable extrapolation technique the wavelengths of corresponding lines in nickel are predicted, which leads to some new identifications and a total of 6 identified lines of Ni xi and 4 of Ni xii.     

On the identification of Ar x and Ar xiv in the solar corona and the origin of the unidentified coronal lines

A study of theZ-dependence of the² P intervals of 2s ²2p and 2s ²2p ⁵, aided by recent observational results, confirms the identification in the coronal spectrum of λ 4412 with Arxiv, and of λ 5533.4 with Arx. It is further shown that transitions from metastable levels in the configurations 3s ²3p ^k 3d, withk=3, 4 and 5, of Fexi,x, ix, and Nixiii, xiii, xi can well account for the remaining unidentified coronal lines.     

A comparison of theoretical line strengths for the 2s 22p 2 − 2s2p 3 transitions in Ne v with solar data

Results are presented for several theoretical line ratios in Nev involving transitions between multiplets in the 2s ²2p ² and 2s2p ³ configurations. A comparison of these with solar data from the S082A and S-055 instruments on board Skylab reveals generally good agreement between theory and experiment, especially in the case of the high-resolution (S082A) observations. However the 2s ²2p ^{2 1} D – 2s2p ^{3 1} P (365.6 Å) and 2s ²2p ^{2 3}P – 2s2p ^{3 3} S (359 Å) lines appear to be blended, possibly with transitions in Fex and Fexi/Fexiii, respectively. We note that the intensity ratio I(365.6 Å)/I(416.2 Å) should be a valuable calibration check for a high-resolution extreme ultraviolet instrument in the spectral range 360–420 Å.     

The identification of new forbidden coronal lines in the solar EUV spectrum

Identifications are proposed for twenty of the twenty-eight coronal lines observed in the spectra obtained during a rocket flight into the path of the 7 March, 1970 solar eclipse. The methods by which the lines have been identified are discussed. Most of the lines identified are from forbidden transitions between levels in the ground 2p ⁿ and 3p ⁿ configurations in high ions of magnesium, silicon, sulphur, iron, and nickel. The temperature range represented is from 6.9 × 10⁵ K to 2.5 × 10⁶ K. The classification of three lines of Fexii and two of Nixiv has led to a revised identification for the near ultraviolet ² D _3/2-² P _1/2 transition in Fe xii. This transition can be identified with the line at 3072 Å rather than that at 3021 Å as previously suggested in the literature.     

Observations of the infrared FeXIII lines in the solar corona of 12 November, 1966

The purpose of this paper is to report on some intensity measurements of the Fe xiii lines at 10 747 Å and 10 798 Å made during the total eclipse of 12 November, 1966. Infrared spectra were taken of the solar corona at a dispersion of 90 Å per mm, using an RCA image converter and spectrograph aboard the NASA CV 990 aircraft off the coast of southern Brazil. The spectra have been reduced to equivalent width in terms of the coronal continuum and values derived for different points in the corona.The observed equivalent widths of the lines lie in the range 10 to 30 Å for the 10 747 line and 5 to 12 Å for the 10 798 line. The ratio of these equivalent widths is found to vary from 2.3 in the inner corona to 6 at a point 1.36 solar radii from the center of the Sun.The above results are discussed in terms of the excitation mechanisms involved in producing the lines. In particular, the results are compared with the recent theoretical calculations of Chevalier and Lambert, who are the first to include the effects of proton collisions in the excitation of the 3p ^{2 3} P levels of Fexiii. Our observations are consistent with an electron density of 4 × 10⁸ in the inner corona; a value which compares favorably with those derived by other observers from the strength of the K continuum. These are, to our knowledge, the first eclipse observations of the infrared Fe xiii lines which indicate that proton collisions are important in the excitation of the coronal lines. The coronal abundance of iron is estimated from the equivalent width of the 10 747 line, and in common with other observers we find an overabundance as compared to the photospheric abundance by a factor of 10.     

A comparison of theoreticalCIV emission line strengths with active region observations obtained with the Solar EUV Rocket Telescope and Spectrograph (SERTS)

Theoretical line ratios involving 2s ² S - 3p ² P, 2p ² P - 3s ² S, and 2p ² S - 3d ² D transitions inCiv between 312 and 420 Å are presented. A comparison of these with solar active region observational data obtained during a rocket flight by the Solar EUV Rocket Telescope and Spectrograph (SERTS) reveals good agreement between theory and experiment, with discrepancies that average only 22%. This provides experimental support for the accuracy of the atomic data adopted in the line ratio calculations, and also resolves discrepancies found previously when the theoretical results were compared with solar data from the S082A instrument on boardSkylab. The potential usefulness of theCIV line ratios as electron temperature diagnostics for the solar transition region is briefly discussed.     

Theoretical intensities of Fexiv in the solar EUV spectrum

Equilibrium population of Fexiv levels in coronal conditions was calculated including configurations 3s ²3p, 3s3p ², 3s ²3d, 3p ³, 3s3p3d, 3s ²4s, 3s ²4p, 3s ²4d, 3s ²4f. Relative populations of selected levels are given in Table VII. Figure 1 shows the dependence of relative intensities of the strongest lines on electron density. Certain line ratios can be used for the determination of N _e .E.g., at T=2 × 10⁶ K and with a dilution factor 0.4, the intensity ratio of λ211.3 and λ219.0 changes by a factor of 65 if N _e increases from 10⁷ to 10¹¹ (Table VIII). Cascades from the 3s3p3d and 3p ³ configurations are important in the population of some levels of 3s3p ² (Table VI). A possibility of identification of additional lines in the solar spectrum is indicated. NAS-NRC Resident Research Associate.     

Comparison of computed fluxes for Fex and Fexiv lines with observed values at 1980 eclipse

Fluxes have been computed for Fex (6374 Å) and Fe xiv (5303 Å) lines as a function of solar radii and at various coronal tempratures. The electron density derived from the white light corona during the total solar eclipse of 1980 were used in the computations. Fluxes in adjacent continua have also been computed. The computed ratios of line flux to the square of continuum flux at a coronal temperature of 1.6 × 10⁶ K show a good fit with the observed values for Fex line. Further, radiative excitation seems to dominate over collisional excitation beyond 1.3 solar radius.     

Solar flare X-ray spectra

Results are presented of an investigation of solar flare X-ray spectra in the region 1.70–1.95 Å, obtained aboard the Intercosmos-4 satellite during the maximum of solar activity (October–November, 1970). With the use of 6 high resolution spectra in the region 1.85–1.87 Å the identification of lines due to 18 transitions of 2p 1s type, consisting of the resonance, intercombination and forbidden Fe xxv ion lines and the satellite Fe xxiv lines has been performed. With the use of the recent laboratory data the averaged wavelengths of the lines were obtained confirming the theoretically calculated ones with an accuracy about ± 0.0004 Å. A variable Doppler shift of the Fe xxv resonance lines was observed for the flare of November 16, 1970, which points to hot plasma motions with velocities up to 400 km s^-1.     

Extreme ultraviolet spectra of solar active regions and their analysis

Extreme ultraviolet spectra of several active regions are presented and analyzed. Spectral intensities of 3 active regions observed with the NRL Skylab XUV spectroheliograph (170–630 Å) are derived. From this data density sensitive line ratios of Mg viii, Si x, S xii, Fe ix, Fe x, Fe xi, Fe xii, Fe xiii, Fe xiv, and Fe xv are examined and typically yield, to within a factor of 2, electron pressures of 1 dyne cm^–2 (n _e T = 6 × 10¹⁵ cm^–3 K). The differential emission measure of the brightest 35 × 35 portion of an active region is obtained between 1.4 × 10⁴ K and 5 × 10⁶ K from HCO OSO-VI XUV (280–1370 Å) spectra published by Dupree et al. (1973). Stigmatic EUV spectra (1170–1710 Å) obtained by the NRL High Resolution Telescope and Spectrograph (HRTS) are also presented. Doppler velocities as a function of position along the slit are derived in an active region plage and sunspot. The velocities are based on an absolute wavelength scale derived from neutral chromospheric lines and are accurate to ±2 km s^–1. Downflows at 10⁵ K are found throughout the plage with typical velocities of 10 km s^–1. In the sunspot, downflows are typically 5 to 20 km s^–1 over the umbra and zero over the penumbra. In addition localized 90 and 150 km s^–1 downflows are found in the umbra in the same 1 × 1 resolution elements which contain the lower velocity downflows. Spectral intensities and velocities in a typical plage 1 resolution element are derived. The velocities are greatest ( 10 km s^–1) at 10⁵ K with lower velocities at higher and lower temperatures. The differential emission measure between 1.3 × 10⁴ K and 2 × 10⁶ K is derived and is found to be comparable to that derived from the OSO-VI data. An electron pressure of 1.4 dynes cm^–2 (n _e T = 1.0 × 10¹⁶ cm^–3 K) is determined from pressure sensitive line ratios of Si iii, O iv, and N iv. From the data presented it is shown that convection plays a major role in determining the structure and dynamics of the active region transition zone and corona.     

Transitions between the n = 2 and n = 3 levels of Ne VII in the solar spectrum

Theoretical populations of the 2s3l levels of Ne vii are presented for electron temperatures from 2.5 × 10⁵ K to 4 × 10⁶ K and electron densities from 10⁸ cm^–3 to 10¹² cm^–3. These, in conjunction with intensities of previously observed solar Ne vii lines and wavelengths and intensities observed in the laboratory, are used to identify further Ne vii lines in the solar spectrum. The dependence on temperature of intensity ratios such as I(2s2p ¹ P – 2s3d ¹ D)/I(2s2p ³ P – 2s3d ³ D) is demonstrated and the advantages of the small wavelength separation of such lines for solar electron temperature diagnostics are discussed.     

A Comparison of Theoretical mgvi Emission Line Strengths with Active-Region Observations From Serts

R-matrix calculations of electron impact excitation rates in N-like Mgvi are used to derive theoretical electron-density-sensitive emission line ratios involving 2s ²2p ³–2s2p ⁴transitions in the 269–403 Å wavelength range. A comparison of these with observations of a solar active region, obtained during the 1989 flight of the Solar EUV Rocket Telescope and Spectrograph (SERTS), reveals good agreement between theory and observation for the 2s ²2p ^{3 4} S–2s2p ^{4 4} Ptransitions at 399.28, 400.67, and 403.30 Å, and the 2s ²2p ^{3 2} P–2s2p ^{4 2} Dlines at 387.77 and 387.97 Å. However, intensities for the other lines attributed to Mgvi in this spectrum by various authors do not match the present theoretical predictions. We argue that these discrepancies are not due to errors in the adopted atomic data, as previously suggested, but rather to observational uncertainties or mis-identifications. Some of the features previously identified as Mgvi lines in the SERTS spectrum, such as 291.36 and 293.15 Å, are judged to be noise, while others (including 349.16 Å) appear to be blended.     

Fabry-Pérot line profiles in the λ5303 å and λ6374 å coronal lines obtained during the 1983 Indonesian eclipse

During the total solar eclipse of 11 June, 1983, an imaging dual-channel Fabry-Pérot interferometer was used to obtain line profiles simultaneously in the green 5303 Å [Fe xiv] and the red 6374 Å [Fe x] coronal lines at various positions in the corona. Extensive microdensitometry followed by multi-Gaussian curve-fitting analysis has resulted in the determination of coronal temperatures and velocity separations between different pockets of coronal gas in the line of sight over a large extent of the corona. Fewer high temperature zones are to be found in the corona of 1983 compared with our similar green-line measurements of the solar maximum corona of 1980. The data are consistent with a temperature maximum occurring at 1.2 R , as found at the 1980 eclipse, but our new data are insufficient to observe farther out than this radius and so determine the position of a maximum. The velocity field in the corona at the 1983 eclipse is less structured compared with that at the 1980 eclipse and is mainly confined to the zone 20–30km s^–1.     

Predicted wavelengths of coronal transitions in the configurations 3s 23p 2, 3s 23p 3 and 3s 23p 4

The level intervals in the ground configurations 3s ²3p ^k(k = 2, 3, 4) have been studied through the isoelectronic sequences up to Ni. Semiempirical expressions for the parameters F ₂(pp) and _p, are derived from observed data and the intervals are recalculated from interpolated or extrapolated values of the parameters. As a result, predicted wavelengths of magnetic-dipole transitions within the configurations are obtained with an estimated uncertainty of one or two Å.     

Spectroscopic Studies of Solar Corona VII. Formation of a Coronal Loop by Evaporation

We obtained time-sequence spectroscopic observations in (Fe x) 6374 Å and (Fe xiv) 5303 Å lines successively with the 25-cm coronagraph, and narrow-band and Doppler images in 5303 Å line by the 2-D 10-cm Doppler coronagraph NOGIS at the Norikura Solar Observatory, of a coronal region for about 7 h on 9 19–20, 2001. The raster scans were obtained with a quasi-periodicity of about 14 min and NOGIS obtained the images with an interval of about 1 min. The coronal region observed showed the formation of a coronal loop by a high-speed surge in the 6374 Å line rising from one of the footpoints of the loop. Off the limb spectroscopic observations in the 6374 Å line showed large velocities along the line of sight and vertical to the solar limb at the time of formation of the loop. The 5303 Å line observations showed negligible line-of-sight velocities and low vertical velocities when compared to those in the 6374 Å line. A hump in the intensity plots in 5303 Å with height appears to move up with respect to the solar limb with an average velocity of 4km s^–1. The FWHM of the 6374 Å showed a much smaller value of about 0.7 Å near the foot point as compared to a value of 1.2 Å at larger heights at the beginning of observations. Later as the loop developed, the FWHM of 6374 Å line showed a gradual decrease along the loop up to 70 from the limb, reached a minimum value of about 0.5 Å and then increased with height during the formation of the loop; this trend lasted for about 2 h. About 3 h after the beginning of the formation of the loop, the FWHM of 6374 Å emission line showed normal values and normal rate of increase with height with some fluctuations. The FWHM of the 5303 Å line did not show such variations along the loop and showed normal decrease in FWHM with height found earlier (Singh et al., 2003a). These observations suggest that a relatively cooler plasma at a temperature of about 0.7 MK or less (corresponding to minimum value of FWHM of 0.5 Å) was ejected from the transition region with a large velocity of about 48km s^–1, heated up in the corona by some process and formed a coronal loop with a height of about 200 above the limb that had lifetime greater than 4 h. It appears that the plasma moved from one of the footpoints and the loop was formed by evaporation of chromospheric plasma. No large-scale brightening and H flare were observed in this region during the observational period of 7 h.On leave from Indian Institute of Astrophysics, Bangalore 560034, India.     

Oscillator strengths for optically allowed transitions in carbon-like ions

Using configuration interaction wavefunctions, the excitation thresholds for the twelve lowest terms arising from the configurations 1s ²2s ²2p ², 1s ²2s2p ³, and 1s ²2p ⁴ of Mgvii have been calculated. The same wavefunctions have been used to calculate the oscillator strengths for the optically allowed transitions in Mgvii. Combining these results with earlier published values for Oiii, Nev, Siix, Caxv, and Fexxi, oscillator strengths for other ions in the carbon isoelectronic sequence (F through Ni) have been predicted. The predicted values are found to be slightly lower when compared with the available published results.     

Multiple loop activations and continuous energy release in the solar flare of June 15, 1973

The spatial and temporal evolution of the high temperature plasma in the flare of 1973 June 15 has been studied using the flare images photographed by the NRL XUV spectroheliograph on Skylab.The overall event involves the successive activations of a number of different loops and arches bridging the magnetic neutral line. The spatial shifts and brightenings observed in the Fe xxiii–xxiv lines are interpreted as the activation of new structures. These continued for four or five minutes after the end of the microwave burst phase, implying additional energy-release unrelated to the nonthermal phase of the flare. A shear component observed in the coronal magnetic field may be a factor in the storage and release of the flare energy.The observed Fe xxiii–xxiv intensities define a post-burst heating phase during which the temperature remained approximately constant at 13 × 10⁶ K while the Fe xxiv intensity and 0–3 Å flux rose to peak values. This phase coincided with the activation of the densest structure (N _e = 2 × 10¹¹ cm^–3). Heating of higher loops continued into the decay phase, even as the overall temperature and flux declined with the fading of the lower Fe xxiv arches.The observed morphology of individual flaring arches is consistent with the idea of energy release at altitude in the arch (coincident with a bright, energetic core in the Fe xxiv image) and energy flow downward into the ribbons. The Doppler velocity of the Fe xxi 1354 Å line is less than 5 km s^–1, indicating that the hot plasma region is stationary.The relation of this flare to the larger class of flares associated with filament eruptions and emerging magnetic flux is discussed.     

X-rays spectroheliograms in lines of Mg xi and Mg xii

Spectroheliograms in the L Mg xii line and in the Mg xi resonance (R) 1s ²¹ S ₀-1s2p ¹ P ₁ line, intercombination (I) 1s ²¹ S ₀-1s2p ³ P _1,2, line, and the forbidden (F) 1s ²¹ S ₀-1s2s ³ S ₁ line, have been obtained.Two Bragg crystal spectrometers were used mounted with mechanical collimators to obtain a spatial resolution of 1 × 3. The apparatus was launched on a sounding rocket on July 2nd, 1971. A particularly thorough study was made of the brightest active region (MC 11402).Variations in the F to I Mg xi line intensity ratio from one point to another in the active region did not reveal the presence of high electron densities.The observed intensities of the Mg xi R line, Mg xii L line and Mg x 1s ²2s ² S _1/2-1s2p ¹ P 2s × × ² P1_{1/2, 3/2} S line are not well explained by an isothermal model. Good agreement between computed and observed intensities is obtained using the non-isothermal model proposed here.     

Line profile analysis of a coronal formation observed near a quiescent prominence: Intensities,temperatures and velocity fields

A technique developed for analysing line profiles with both speed and high accuracy was used to study the physical conditions of a coronal formation near a quiescent prominence. Detailed analyses of five coronal lines (Fe xiv λ 5303, Fe x λ 6374, Ni xv λ 6702, Fe xv λ 7059, and Fe xi λ 7892) provided total intensities, Doppler width temperatures, ionization temperatures, and velocities. Dissimilar spatial fluctuations in intensity are obvious for ions grouped according to (low vs high) ionization potentials. The intensity of the green line shows a local minimum around the observed quiescent prominence; a corresponding but much more diffuse pattern is visible in the red line intensity. Large differences are observed in temperatures derived by different means. In particular, , while , and . The differences between and are taken as direct evidence of temperature inhomogeneity. One can thus put little significance in T _e (xi/x). T _D(λ5303) and T _e (xv/xiv) fluctuate nearly in parallel at each slit height, with a weak local minimum evident around the prominence. The discrepancy between these two can be removed if a non-thermal turbulent motion of 6–16 km s⁻¹ is assumed. Variations with height of both T _D(λ5303) and T _e (xv/xiv) suggest that the coronal temperature maximum is located no more than 15000 km above the top of spicules. A negative gradient of about 6 deg km⁻¹ is found in the height variation of T _D(λ5303). The height variation of the green line wavelength shows that the majority of coronal material in this region is flowing from west to east on the Sun, with the highest velocity of 12 km s⁻¹ found at the lowest heights. This motion is in the same sense as that of the nearby coronal rain, as determined both from the spectra and wavelength-shifted Hα filtergrams. Superposed on the above flow is a systematic velocity field of up to ±5 km s⁻¹. This field similarly reaches maximum amplitudes at lowest heights showing a local maximum around the prominence. On leave from Institute of Earth Science and Astrophysics, Shiga University, Ohtsu 520, Japan, as 1973–75 National Academy of Science/National Research Council Senior Post-Doctoral Research Associate at Sacramento Peak Observatory.