NEON AND OXYGEN ABSOLUTE ABUNDANCES IN THE SOLAR CORONA

The Astrophysical Journal, 659:743Y749, 2007 April 10 # 2007. The American Astronomical Society. All rights reserved. Printed in U.S.A. NEON AND OXYGEN ABSOLUTE ABUNDANCES IN THE SOLAR CORONA E. Landi and U. Feldman Artep, Inc., Naval Research Laboratory, Washington DC 20375-5320 and G. A. Doschek Naval Research Laboratory, Washington DC 20375-5320 Received 2006 October 16; accepted 2006 December 14 ABSTRACT In the present work we use the UV spectrum of a solar flare observed with SOHO SUMER to measure the absolute abundance of Ne in the solar atmosphere. The measurement is carried out using the intensity ratio between the allowed 1s2s 3 S 1 Y1s2p 3 P 2 Ne ix line at 1248.28 8 and the free-free continuum radiation observed close to the Ne ix line. We find a value of the absolute Ne abundance A Ne ¼ 8:11 0:12, in agreement with previous estimates but substantially higher than the very recent estimate by Asplund et al. based on the oxygen photospheric abundance and the Ne/O relative abundance. Considering our measured A Ne value, we argue that the absolute oxygen abundance of Asplund et al. is too low by a factor 1.9. This result has important consequences for models of the solar interior based on helioseismology measurements, as well as on the FIP bias determination of the solar upper atmosphere, solar wind, and solar energetic particles. Subject headinggs: Sun: corona Sun: flares 1. INTRODUCTION The measurement of element abundances is of crucial importance in solar physics, as the chemical composition plays an important role in solar models. For example, element abundances are very important in building models of the solar interior, as they play a key role in determining the opacity of the solar plasma; the chemical composition of different regions of the solar upper atmosphere can be used as tracers of the origin of the solar wind and of flaring plasma. Element abundances are crucial for determining the radiative losses in the solar atmosphere and, hence, play a key role in the energy budget of theoretical models attempting to explain most physical processes taking place in the solar chromosphere, transition region, and corona. It is no surprise that the measurement of the solar element composition has been a major point of interest for solar physicists. The 15 most abundant elements in the solar photosphere are composed of two distinct groups: volatile and nonvolatile elements. The abundances of nonvolatile elements Na, Mg, Al, Si, Ca, Fe, and Ni were measured in the laboratory from C i carbonaceous chondrite meteorites yielding accuracies down to 5%Y10% or better. The abundances of volatile elements are more difficult to measure, since they are not retained in meteorites and must be determined by other means. The solar photospheric spectrum has been the most important source of element abundance measurement for H, C, N, and O, since the spectral lines that they emit when they are not ionized have relatively low excitation energies and, therefore, are present in the solar photospheric spectrum. The abundances of He, Ne, and Ar instead could not be determined in this way, since the excitation energy of lines emitted when these elements are in the neutral state is too high. The absolute abundance of Ar was only recently directly measured from flare spectra using intensity ratios of He-like Ar xvii to free-free emission ( Phillips et al. 2003). However, in many cases, measurements of the Ne abundance relative to O or to nonvolatile elements Mg and Fe from spectra of the solar upper atmosphere and from solar energetic particle events were used to determine 743 the absolute abundances of Ne. Reviews of such measurements are given by Drake & Testa (2005, see their supplementary data) and Acton et al. (1975). Among all elements, C, N, O, and Ne have a special importance, since they are large contributors to the solar metallicity and to the opacity in the solar interior and are therefore key players in models of the inner layers of the Sun. The set of abundances available until recently for these ions, based mostly on solar photospheric spectra for C, N, and O, and on the assumption that the relative Ne/O abundance is Ne/O ¼ 0:15, has allowed modelers of the solar interior to develop predictions of the inner structure of the Sun in excellent agreement with helioseismological data. However, a new generation of three-dimensional non-lte models of the solar photosphere (Asplund et al. 2000; Allende Prieto et al. 2001) led to a recalculation of the atomic level populations that are used to compute the intensities and profiles of spectral lines used to measure element abundances from photospheric spectra and of their blends. These models have led to a substantial decrease of the O abundance (Asplund et al. 2004, 2005) that has had deep impacts on the calculation of the opacity of the solar interior leading to disagreements between models and helioseismology observations. Assuming that the revised oxygen abundance is correct, the disagreement can be cured if the absolute Ne abundance is increased to values in the range of 8:24 logðne/hþ 8:44 (Antia & Basu 2005; later recalculated as 8:29 0:05 by Bahcall et al. [2005]), thus significantly increasing the Ne/O abundance ratio of 0.15 measured in the Sun and in other astrophysical objects. A considerable controversy has flared about the value of the Ne/O abundance ratio, as Drake & Testa (2005) reviewed the value of the Ne/O ratio in nearby stars, finding it to be 0.41, thus bringing the Ne absolute abundance to 8.29 and reconciling the models of the solar interior with those of the photospheric layers. However, recent measurements of the Ne/O abundance in the solar atmosphere obtained from quiet Sun spectra by Young (2005) have confirmed the Ne/O ¼ 0:15 abundance ratio, while active region measurements from Schmelz et al. (2005) provided upper limits to the Ne/O abundance ratio consistent with

744 LANDI, FELDMAN, & DOSCHEK Vol. 659 the standard 0.15 value. These results questioned the enhanced Ne abundance as a solution to this controversy. Clearly, it is very important to obtain measurements of the absolute abundances of Ne and O in the Sun in order to validate or dispute the predictions of the models of the solar photosphere that have sparked this controversy and determine the abundances of these two key elements with sufficient accuracy to enable modelers of the solar interior and of the solar photosphere to predict the structure of the solar interior and photospheric emission with precision. The aim of this paper is to directly measure the absolute abundance of Ne by using a flare spectrum in the ultraviolet (UV) spectral range, by means of the intensity ratio between one Helike Ne ix line and free-free emission. Unfortunately, no lines are available in the UV spectra considered by us that allow us to measure the O abundance. However, using our measured absolute Ne abundance and the observed Ne/O ratio, we also determine the absolute abundance of O. The diagnostic technique and atomic data used to measure the absolute Ne abundance are described in x 2. The observations we have analyzed are outlined in x 3, and the measurement results are reported in x 4 and summarized and discussed in x 5. 2. DIAGNOSTIC METHOD AND TOOLS The diagnostic technique we used to measure the absolute abundance of Ne consists of the comparison of the observed ratio between the He-like Ne ix line observed by the SUMER instrument on board the Solar and Heliospheric Observatory (SOHO) at 1248.28 8 with the free-free continuum observed very close in wavelength. Given the high temperature and emission measure required to observe the free-free continuum in the UV, this measurement has been carried out using the spectrum of a medium-size (M7.6) flare observed by SUMER on 1999 May 9. The use of the free-free emission in place of line intensity ratios between the Ne ix line and the H i lines, also present in the flare spectrum we used, allows us to minimize the effects of uncertainties in the plasma temperature, since the free-free emission is slowly varying with temperature, while H i emission is strongly temperature dependent (as shown in Fig. 1). In addition, contamination of H i line intensities from instrument-scattered light emitted by the disk is avoided. The atomic data used for the calculation of the theoretical ratios to be compared with observations mostly come from version 5.2 of the CHIANTI database (Landi et al. 2006). However, we have slightly modified the calculation of the Ne ix line intensities in order to use more recent data for recombination into excited levels in the 1s2s and 1s2p configurations and to relax the coronal model assumption underlying the CHIANTI calculation in order to explicitly include recombination into the level population calculation, following the approach outlined by Gabriel & Jordan (1969, 1973). The data for recombination into excited configurations have been taken from Porquet & Dubau (2000) and replace the older Mewe et al. (1985) rate in CHIANTI 5.2. This approach and the new recombination data will be adopted in the next version of CHIANTI. Their effects on the emissivity of the Ne ix 1248.28 8 line are moderate (within 10%Y20%) and are due to the change in the recombination rates, since the density of the flare we consider in the present work is low enough to keep the coronal model assumption valid. The free-free continuum and Ne ix line emissivity have been calculated using the Mazzotta et al. (1998) ion abundances. For the calculation of the free-free continuum emissivities, the extended coronal abundances from the CHIANTI database have been used; original values come from Feldman (1992), Landi et al. (2002), Fig. 1. Emissivities of the Ne ix, H i, and free-free emission studied in the present work. Free-free emissivity is reported at 972.5, 1025.7, and 1248.3 8, corresponding to the wavelengths of the spectral lines considered in the present work. and the photospheric values of Grevesse & Sauval (1998) enhanced by a factor of 3.5 to reflect the standard composition of the active solar corona. The use of a different set of abundances has very limited effects (a few percent) on the free-free emissivities. The emissivity of the Ne ix line has been calculated assuming unity abundance, so that the value of the latter can be directly determined from the line-to-continuum ratio. In order to measure the plasma temperature, we also made use of the intensities and line widths of the H i 972.5 and 1025.7 8 lines after scattered light subtraction. The emissivities of the H i lines have been calculated using the method and atomic data described in Feldman et al. (2005) and Laming & Feldman (2001), where the reader is referred for further details; calculations included full treatment of electron impact ionization, proton and electron collision excitation rates. 3. INSTRUMENT AND OBSERVATIONS SUMER is a stigmatic slit spectrometer with a spatial resolution of about 1 00 pixel 1 along the slit length and a spectral resolution of about 43 m8 pixel 1 in first order and about 22 m8 pixel 1 in second order. The wavelength range accessible with the detector used to carry out the present observations (detector A) is 780Y1600 8 in first order and 390Y800 8 in second order. Because of the rapid falloff of sensitivity below 600 8, the real lower wavelength limit for second-order lines is near 500 8. Spectra are recorded by the SUMER detector in 43 8 wide sections. Details of the SUMER instrument can be found in Wilhelm et al. (1995). On 1999 May 9 the SUMER spectrometer was pointed at an active region above the northwest solar limb and was recording a complete SUMER spectrum with the 1 00 ; 300 00 slit and detector A when a flare erupted in the field of view at 18:00 UT, with SUMER recording the 43 8 wide spectral section around 1440 8. SUMER observed the flare continuously for the following 4 hr, recording the remaining 1440Y1600 8 range and then another complete spectral scan (780Y1600 8) of the flare plasma as the flare went through the decay phase. The exposure time for each spectral window was 300 s. The flare was studied in previous works by Innes et al. (2001), who investigated the initiation of a partial-halo coronal mass ejection triggered by the flare, and Landi et al. (2003), who measured

No. 1, 2007 Ne ABSOLUTE ABUNDANCE 745 the main physical parameters of the flare plasma in locations where the emission from the hottest ions in the data set ( Fe xxii, Fe xxiii) was strongest. In the present work we consider the region of the slit where free-free and Ne ix emission are strongest, and Fe xxii and Fe xxiii are minimum, so the diagnostic results obtained by Landi et al. (2003) cannot be applied to the present work. The spectra emitted during the 1999 May 9 flare were destretched and corrected for the flat-field calibration by using the standard SUMER software. The intensity of the Ne ix line, of the free-free continuum, and of the lines required to carry out plasma diagnostics were measured by summing the counts under the line profile after subtraction of the background along the spectral direction or, in the case of the continuum, measuring the total counts in a spectral region a few tens of pixels wide where no spectral line was present. In the case of H i lines at 972.54 and 1025.72 8 a second set of measurements was carried out by fitting a Gaussian line profile to the lines after the background along the slit direction was subtracted (see x 4), so that the value of the line width could also be determined. In each SUMER frame, there was at least one region in the detector where no lines were present and the measurement of continuum intensity was possible. Background subtraction is described in more detail in x 4. 4. MEASUREMENT OF THE NEON AND OXYGEN ABSOLUTE ABUNDANCES 4.1. Emitting Region Selection For the present study, we need to select an emitting region whose physical properties are as homogeneous as possible; in particular, it is very important that the electron temperature be measured with accuracy and that the plasma be close to isothermal, since the Ne ix line emissivity is strongly temperature sensitive outside the 6:4 log T e 6:7 range (T in K). The plasma in the field of view of the SUMER slit included several different components whose position along the slit changed with time during the observations. The Ne ix line and free-free continuum have been observed both before and after the flare onset, but the plasma before the flare onset was considerably complex, being composed of several overlapping structures; in this situation, selecting a nearly isothermal region was impossible. Moreover, the choice of the surrounding ambient emission (the ambient background) to be subtracted from both the Ne ix and the freefree emission was very subjective and had large effects on the resulting abundance measurement. In addition, the presence of cold material dominating the H i emission in the same position of the slit where Ne ix was observed further complicated the use of H i lines and the continuum for temperature measurement. All these combined problems caused the measured absolute Ne abundance to vary by 1 order of magnitude according to the different ambient background or subregion chosen, so the data set before the flare onset is not further considered. The plasma after the flare onset was also strongly multithermal along the slit, but the region emitting the Ne ix and free-free emission was bright and clearly outlined, and the subtraction of the ambient background was relatively easy. Figure 2 shows the image of the slit at several different wavelengths, corresponding to lines of ions having their maximum fractional abundance in the 10 4 Y10 7 K temperature range. Each line is isolated and the ambient background can be easily subtracted. Figures 3 and 4 show the profile of the intensity along the slit length for several lines in the data set. Figure 3 shows the Ne ix line intensity profile having a pronounced maximum between pixels 165 and 230; and this behavior is shared by a number of lines from different ions formed at temperatures between 3 and Fig. 2. Image of the 300 pixel long SUMER slit in several lines emitted at different temperatures. Horizontal lines correspond to pixels 165 and 230. 5MK(6:5 log T e 6:7), some of which are shown in Figure 3. The similarity between the intensity behavior along the slit between pixels 165 and 230 is a clear indication that the same plasma emitted the free-free continuum, the H i lines, and the coronal lines. Therefore, both the free-free emission and the H i line intensities can be used for measuring the physical properties of the plasma emitting the Ne ix line. Several other lines in the present data set do not share the same intensity profile along the slit shown in Figure 3; a few examples are displayed in Figure 4. These lines are emitted by ions with maximum fractional abundance outside the 6:5 log T e 6:7 range, where the lines shown in Figure 3 have maximum emission. Their different intensity profiles indicate that their emission between pixels 165 and 230 is strongly contaminated by plasma from many different structures along the line of sight. On the contrary, lines shown in Figure 3, as well as all lines sharing that same intensity profile, are emitted by the same plasma. In the present work, we have selected the region where the emission of Ne ix,hi, and free-free emission is maximum in order to minimize uncertainties in the measurement of the Ne absolute abundance. This was done for a number of reasons.

746 LANDI, FELDMAN, & DOSCHEK Vol. 659 Fig. 3. Examples of intensity profiles of H i 1025.7 8 and of lines formed at temperatures similar to the peak emission temperature of Ne ix. First, the signal-to-noise ratio for these three spectral features is maximum, and the region can be easily separated from the ambient background emission (see x 4.1). Second, in the 6:5 log T e 6:7 temperature range, the Ne ion abundance has its maximum, and its dependence on temperature itself is much more limited than outside that temperature range, so the effects of uncertainties in the measured value of T e on the predicted Ne ix line emissivity are minimum. Third, this region allows us to use H i emission for temperature diagnostics; in no other part of the slit is this possible, since the H i line intensity has a very Fig. 4. Examples of intensity profiles of lines formed at temperatures very different from the peak emission temperature of Ne ix.

No. 1, 2007 Ne ABSOLUTE ABUNDANCE 747 different profile than that of any other ion. In addition, the removal of the ambient background was less ambiguous and subjective for this region than for any other region along the slit where Ne ix emission was recorded. The use of H i lines for plasma diagnostics is important, since the atomic data necessary to calculate the H i emissivity are more accurate than those available in the literature to calculate the emissivity of all the ions that share the same intensity profile of the Ne ix line. Therefore, by using H i we also minimize the effects of uncertainties in the atomic data on the temperature measurements. It is important to note that recombination continuum in the selected regions is negligible in the present data set, since the free-bound continuum at wavelengths above 912 8 is due to recombination to neutral C, Si, and S. In fact, the abundance of recombining ions is completely negligible at temperatures typical of active regions and flares; the presence of cold loops crossing the selected region can also be dismissed, as they would have caused the intensity profile of the H i line to greatly depart from that shown by Ne ix. Moreover, even in the regions of the slit where cold loops are present (i.e., pixels 280Y290 in Fig. 3) no trace of free-bound continuum emission is found at wavelengths larger than 912 8 both before and after the onset of the flare. 4.2. Background Subtraction Background subtraction consists of two steps. The first step consists of removing the background continuum emission along the spectral direction under the line profile. The second and more delicate step is to remove the ambient continuum emission and the instrument-scattered light from the local emission of the flare plasma. The ambient continuum emission and the instrument-scattered light were removed after the line intensity was measured, summing all counts under the line profile and removing the background along the spectral direction by using the emission at the sides of the spectral line devoid of the presence of any other spectral line. The ambient background was subtracted from the measured line intensity by using the intensity profile of the line along the slit. Figure 5 shows an example, displaying the intensity of the Ne ix 1248.28 8 line along the slit length. In Figure 5 the subtraction of the ambient background intensity from the selected region, marked by the vertical bars, is not uniquely determined and is left with a residual uncertainty in the choice of the portions of the slit to be considered for the calculation of the background, approximated there with a linear fit. However, the strength of the peak of the selected region causes the uncertainty in the background selection to change the Ne ix line flux by 18% at maximum. The uncertainties in the background subtraction for H i and the free-free emission are 20%. A different procedure was adopted for removing the background from the H i line profile, necessary to measure the H i line width. In this case, the ambient background was removed first for each pixel along the spectral direction resulting in a cleaned spectrum. The H i line profile was then fitted with a Gaussian shape and a linear background along the spectral direction; the line full width at half-maximum (FWHM) was then determined from the fitted line width. 4.3. Temperature and Density Measurement The highly ionized ions whose intensity profile is similar to that of Ne ix are: Ar xii, xiii, Caxiv, xv, Tixv, xvi, Crxvii, xviii, Mn xvii, and Fe xvii, xviii. With the only exception of Fe xvii, xviii, atomic data for these ions are inadequate to measure the temperature structure of the emitting plasma because of the limited size of the atomic model (Ti, Cr, Mn, and Ar xiv) that does Fig. 5. Ambient background subtraction from the intensity profile of the Ne ix 1248.28 8 line along the slit. The difference in the final intensity of the Ne ix line is 18%. not allow one to take into account the contribution of cascades from high-energy levels and because of the lack of resonances in the collision data (Ar xii,caxiv). These limitations of the data in CHIANTI and in the literature and the lack of other lines from other ions with more accurate atomic data prevent an accurate measurement of the temperature structure of the emitting plasma both with line intensity ratios and with a differential emission measure analysis. In order to avoid all these atomic physics problems, we have used the H i lines and free-free continuum intensities to measure the temperature of the emitting plasma. The electron density is of limited importance, as the free-free, H i, and Ne ix emissivities are largely independent of it in the range of values typical for active regions and moderate flares; however, Ar xii, xiii and Ca xiv line ratios consistently indicate a density of log N e ¼ 10:1 0:2(N e in cm 3 ), which is used to calculate the line and continuum emissivities. The plasma temperature has been measured using two independent techniques: from the intensity ratio of the H i 972.5 and 1025.7 8 lines with the free-free intensity and from the line width of the two H i lines. The temperature values provided by these two diagnostic techniques are in good agreement. Figure 6 displays the calculated and measured intensity ratios between the two H i lines and the free-free continuum, showing that the two lines provide temperature values between log T ¼ 6:55 and 6.60 (T in K). In order to determine the plasma electron temperature from the H i line widths, we used the standard formula dw ¼ k 0 c sffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi 4ln2 2kT H i M þ v nth 2 ; ð1þ where dw is the FWHM, v nth is the nonthermal velocity, T H i is the H i ion temperature, k is the Boltzmann constant, M is the hydrogen mass, c is the speed of light, and k 0 is the centroid wavelength. The use of equation (1) in the present work requires three assumptions: 1. that the instrumental width is negligible or corrected before the measurement is carried out, 2. that the plasma nonthermal velocities are negligible or measured by some other technique, 3. and that the electron and H i ion temperatures are the same.

748 LANDI, FELDMAN, & DOSCHEK Vol. 659 regions, we can determine the absolute abundance of oxygen to be A O ¼ 8:93 0:12, in agreement with Grevesse & Sauval (1998), but higher by a factor of 1.9 than Asplund et al. (2005). Unfortunately, no direct measurements of the O absolute abundance were possible in the present data set because of the strong temperature dependence of the O vi emissivities in the log T ¼ 6:61 0:06 temperature range and the difficulties of subtracting the ambient background emission along the slit to the measured intensity of these lines. Moreover, He-like O vii lines are found between 1620 and 1640 8, just outside the SUMER wavelength range; these lines would have proved very useful in estimating the O absolute abundance with the same technique used in the present paper. The lack of a direct measurement of A O does not allow us to confirm unambiguously the solar absolute oxygen abundance A O ¼ 8:93 0:12 that we propose. Fig. 6. Calculated intensity ratios between free-free and H i spectral lines. Measured ratios are also indicated. In the case of H i lines, the instrumental width can be neglected, because the small atomic weight of the hydrogen nucleus ensures that random thermal velocities are so high to broaden a H i line well beyond any instrumental effect (FWHM 1:2 8 for the two lines we are considering). The large thermal velocities also minimize the effects that the nonthermal velocities have on the line width. In order to take into account v nth,wehavemeasuredthe temperature assuming three values of v nth, 20, 50, and 80 km s 1, typical offlare decay phases and found that the change in the measured temperature is log T ¼ 0:06. The third assumption is verified in plasmas with densities of order 10 10 cm 3,suchasthe present one, since the plasma equilibration times are shorter than the timescales of the flare decay phase we are studying. The values of the temperatures given by the two lines are the same within uncertainties, and we take them to be log T ¼ 6:61 0:06 (T in K). The values of the measured temperatures correspond to the temperature of maximum emission of the Ne ix 1248.28 8 line, so uncertainties in the emissivities due to the temperature are reduced to 15% at maximum. 4.4. Neon Absolute Abundance The free-free/ Ne ix intensity ratio observed from the flare was compared to the theoretical value calculated using the electron temperature and density values measured from the observed spectra, assuming unity Ne abundance. The abundance A Ne of Ne relative to hydrogen can then be determined as ð A Ne ¼ log I A=I 1248 Þ pred ði A =I 1248 Þ obs þ12; ð2þ where the hydrogen number density has been assumed to be log N(H)¼ 12, as it is usually done in the literature. The uncertainties in the theoretical ratio are due to the variations of the Ne ix emissivity due to the uncertainties in the temperature, while those in the free-free emissivity are negligible due to the moderate dependence of the latter on T e. Uncertainties in the observed ratio are driven by the uncertainties in the background subtraction to both the free-free and the Ne ix 1248.28 8 line. Considering these uncertainties, the resulting Ne absolute abundance is A Ne ¼ 8:11 0:12. 4.5. Oxygen Absolute Abundance If we also assume that Ne/O ¼ 0:15, as confirmed by Young (2005) and Schmelz et al. (2005) for quiet-sun and cool active 5. SUMMARY In the present work we have used the intensity ratio between a Ne ix line and the free-free continuum intensity measured during the decay phase of a flare with SUMER to measure the neon absolute abundance of A Ne. Using the Ne/O relative abundance measured by several authors in the solar upper atmosphere, we also determine the oxygen absolute abundance A O. The values that we find are A Ne ¼ 8:11 0:12 and A O ¼ 8:93 0:12. The present measurement of the absolute abundance of Ne is the first done from the far-uv spectrum of the Sun and provides for the first time a direct measurement of A Ne in the solar atmosphere, without the need of a reference element, as in previous estimates of the Ne abundance (Feldman & Widing 2007). This value is in agreement with earlier estimates from Grevesse & Sauval (1998), with solar energetic particle measurements (Reames 1999), and with the spectroscopic measurements in the solar upper atmosphere reviewed by Feldman & Widing (2007); however, it is higher than the value proposed by Asplund et al. (2005) by a factor of 1.9. The value from Asplund et al. (2005) was determined from the absolute abundance of oxygen they provide and the widely accepted Ne/O abundance ratio of 0.15. It is to be noted, however, that past measurements of the Ne/O abundance ratio summarized by Drake & Testa (2005) and Acton et al. (1975) show considerable scatter, even among measurements of different active regions and flares made with the same instrumentation. This scatter is well beyond the uncertainties provided by each of the works reported there. Such scatter might suggest that the Ne/O ratio is not constant in the solar atmosphere, especially in active regions and flares, possibly because of some fractionation process(es) operating on one or both elements. If such processes were active in the plasma we considered in the present work, there would be two consequences for our results: (1) the oxygen absolute abundance that we derived from the Ne/O abundance ratio of 0.15 might not be correct, and (2) the absolute neon abundance we measured might not be representative of the real value in the Sun. We recommend further studies of the Ne/O abundance of the Sun be carried out to investigate the Ne/O relative abundance variability. The work of Enrico Landi is supported by NASA grants. The work of Uri Feldman and George Doschek is supported by NASA grants and NRL 6.1 funding. We thank the referee for his valuable comments that helped improve the original manuscript. CHIANTI is a collaborative project involving the NRL ( USA), RAL ( UK), MSSL ( UK), the Universities of Florence ( Italy) and Cambridge ( UK), and George Mason University ( USA).

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