Oxygen Abundance in Polar Coronal Holes
L. Teriaca*, G. Poletto*, A. Falchi* and J. G. Doylef
*Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5,50125 Firenze, Italy
^Armagh Observatory, College Hill, BT61 9DG Armagh, N. Ireland
Abstract.
Fast solar wind is known to emanate from polar coronal holes. However, only recently attention has been
given to the problem of where, within coronal holes, fast wind originates. Information on whether the fast solar
wind originates from plumes or interplume regions may be obtained by comparing the elemental abundances in
these regions with those characterizing the fast wind. Here we present a first attempt to determine the oxygen
abundance in the interplume regions by using spectra taken at times of minimum in the solar cycle (when it is
easier to identify these structures) by the SUMER spectrograph aboard SoHO. To this end, we analyze spectra
taken in 1996 in polar regions, at altitudes ranging between 1.05 and 1.3 R0, finding a value > 8.5 for the oxygen
abundance in the interplume regions. From the analysis of the O VI 1032 to 1037 line intensity ratio we also
find no evidence of outflow velocities below 1.2 solar radii in interplume regions, while there are indications
that outflow motions start to be significant above 1.5 solar radii. The method used and the assumptions made are
discussed in light of the derived values. Our values are compared with previous determinations in the corona and
solar wind.
INTRODUCTION
Ulysses observations have identified a difference in
the elemental abundances of the fast and slow wind: low
FIP elements being more enriched, with respect to their
photospheric abundance, in the slow speed wind than in
the high speed wind. Hence, if interplume regions are
really the sources of the fast wind, we may possibly find
a difference in the elemental composition of plume and
interplume plasma at coronal levels.
The purpose of the present work is to shed some light
on the interplume composition by determining which
oxygen abundance is consistent with the O VI1032,1037
A line intensities measured in the range between 1.05
and 1.3 R0. It is also necessary to point out that the FIP
bias usually refers to the comparison between high FIP
and low FIP elements. This paper deals with absolute
abundance instead, which actually is a more appropriate
parameter in distinguishing the coronal origin of the solar
wind.
Polar coronal holes have long been recognized to be the
sources of the high speed solar wind, but only recently it
has been identified where, within coronal holes, the solar
wind originates. UVCS and SUMER measurements of
line widths and of the ratio of the O VI doublet lines
at 1032 and 1037 A in the low corona (at heliocentric
distances lower than 2.5 RQ) seem to indicate that the
low density background plasma, rather than the high
density plume plasma, is the site where high speed wind
originates. Recent analyses of SUMER and UVCS data
[see e. g. 1, 2, 3] have shown that the width of UV lines
is larger in interplume than in plume regions, hinting
to interplumes as the site where energy is preferentially
deposited and, possibly, fast wind emanates. Moreover,
the analysis of fast polar wind performed by [4] seems
to favor the low density regions as sources of the fast
wind streams, while a work by [5] shows that there is no
evidence of outflow motions in bright points and plumes
observed within a coronal hole in the Ne vm 770 A line.
A further factor that may differentiate plume and interplume regions is their elemental abundance: a FIP bias
(overabundance of particles with low First lonization Potential over photospheric abundances) in plumes of ~ 15
has been found by [6]. Because only small FIP bias of
up to 2 is found in the fast wind [7], we may expect that
interplumes would not show large FIP bias either.
OBSERVATIONS
The observations discussed here were acquired by
SUMER on 3 June 1996 and comprise two rasters of
the North Polar Coronal Hole (NPCH). The first part
is a raster scan consisting of 75 spectra obtained using
the 4" x 300" slit with an exposure time of 60 s and
covering an area of ~ 278" x 300" centered at helio-
CP598, Solar and Galactic Composition, edited by R. F. Wimmer-Schweingruber
© 2001 American Institute of Physics 0-7354-0042-3/017$ 18.00
65
Polar plumes are linear structures that are apparent
over the solar poles in visible light, in extreme ultraviolet
and in soft X-rays [see 8, and references therein]. Despite
the fact that bright plumes are striking features within
coronal holes that can extend to more than 30 solar radii
from the Sun [8, 9], they only account for a tiny fraction
of the solid angle subtended by the polar coronal holes.
Plumes are well visible in O VI lines and our rasters show
at least three of these structures. After selecting a plume
and an interplume region, data within these regions were
binned over 5 x 15 pixels (in the x and y direction,
respectively) for the first (lower) scan, and over 4 x 30
pixels for the second scan. In such a way, line profiles
were determined at 15 different locations above the solar
limb for both plume and interplume (see Fig. 1). Further
details on this dataset can be found in [2].
1.50
1.40
0
K
PS
1.30
DATA ANALYSIS
CO
Spectral profiles were carefully fitted in order to obtain
C II 1037, O VI 1032 and O VI 1037 A line intensities.
Despite the very high quality of the telescope mirror,
the level of stray light is not negligible when observations
of lines that are bright on disk are carried out above the
limb [10, 11].
A chromospheric line, such as C II 1037, can be assumed not to have any coronal contribution and, hence,
to be entirely due to stray light. The O VI1032 stray component at a certain altitude above the limb will be, hence,
given by the product of the C II 1037 intensity at that
altitude times the ratio of the disk-averaged O vi 1032
and C II 1037 line intensities. The above ratio was evaluated using observations at 1.8 RQ, where the SUMER
spectrum is entirely due to stray light [10, 12, 2].
Errors on line intensities were calculated through Poissonian statistics and, finally, their propagation in the O VI
1032/1037 line ratio was established.
1.20
1.10
-100
0
100
Solar X (arcsec)
FIGURE 1. Intensity map of the north polar coronal hole as
seen by SUMER in O VI on June 3rd 1996. Strong plumes
and inter plume lanes can be identified in this diagram out to
1.5 RQ above the limb. The contrast was enhanced by normalizing to the average intensity profile in the j-direction. Areas
enclosed by white (black) solid lines are the ones over which
an integration was performed in order to obtain the the averaged
line intensities in plumes (interplumes) regions discussed in the
text. The averaged areas measure ~ 19" x 15" and ~ 23" x 30"
in the lower and upper dataset, respectively
RESULTS AND DISCUSSION
In the case of allowed transitions from the ground level
of ions without metastable levels (such as Li-like ions),
only the ground level g and the excited upper level u responsible for the line emission are important for calculating the line intensity and the transition is, hence, well
described by a two level atomic model. This is the case
of the O VI 1032 and 1037 A lines. In the solar corona,
these lines are formed by electron impact excitation (collisional component) and by resonant absorption of the
O VI radiation from the transition region (radiative com-
centric coordinates X = —4", Y = 1150". An exposure
time of 150 s and the 4" x 300" slit were used for the
second part, consisting of 52 spectra covering a region
of - 287" x 300" centered at X = -6", Y = 1300"
(see Fig. 1). These rasters were aimed to obtain high
signal-to-noise line profiles in the plume and interplume
regions above the limb out to 1.5 RQ. Data were reduced
using various IDL routines from within the SUMER
software tree.
66
with the assumption of ionization equilibrium (implied
in the calculation ofR(Te)) and with the uncertainties in
the knowledge of the atomic data (see Raymond et al,
present proceedings).
/Q was evaluated assuming a l/cos(G) center-to-limb
line intensity variation ([19]) and adopting an intensity
at disk centre of 280 mW m~ 2 Sr1 for the O vi 1032
A line. The above value was obtained from a disk centre
spectrum obtained on June 4th 1996.
In the case of the observations here discussed, we can
assume the magnetic field lines to be perpendicular to
the line of sight. This allow us to identify T^~ with the
effective temperature Teff defined through the equation:
12
10
O
O
1.00
1.10
1.20
1.30
1.40
Radial Distance in R Q
*?/.=
1.50
(3)
where A is the wavelength, c the velocity of light and AA/)
the measured half width of the line at 1 je of the peak intensity. lRad is not a strong function of T^~ and, hence,
a vi/e value of 65 km s"1 was assumed, for the present
work, as representative of both plume and interplume.
The above value was obtained averaging the results published by [2]. Furthermore, TO has been assumed equal
ponent),
=
x
boll. + lRad. = 0-85
(2)
where m is the ion mass and v\/e is the Doppler width
(in km s"1) or most probable speed of the unresolved
(thermal + non-thermal) motions distribution, given by
FIGURE 2. Electron temperature as a function of altitude
in polar coronal holes. Measurements obtained in interplume
(filled circles) and plume (open circles) regions, respectively
[3]. Values obtained without distinguishing between plume and
interplume: crosses [11]. Dashed and solid lines represent the
adopted temperature profiles in plume and interplume, respectively.
lobs
m
-sgu
4n
[(q(Te)R(Te)N*) + BguIQ(D(To,W)R(Te)Ne)]
toi;.
The determination of the electron temperature profile
in the solar corona is particularly controversial. The discrepancy between the electron temperatures obtained in
the inner corona through spectroscopic diagnostics [see
e. g. 3] and the values deduced from ion charge composition measurements obtained in-situ in the high-speed
solar wind by Ulysses [see e. g. 20], is well known. Spectroscopically derived Te values never exceed 106 K in
the range 1.05-1.3 R0 (see Figure 2). Values by [3] are
particularly accurate since they use the ratio of lines of
the same ion (Mg IX 706/750), thus avoiding problems
with elemental abundances, ionic fractions, and differential flow speeds between different ions [21]. The closeness in wavelength minimize also the errors arising from
the instrument calibration. Moreover, these are the only
measurement of Te in plume and interplume. It has been
shown by [21] that in the region 1.05-1.5 RQ this ratio
is scarcely affected by the presence of a small percentage of electrons following non-Maxwellian distribution
functions. From the above discussion we adopted the Te
profiles shown in Figure 2.
The O VI 1032/1037 line ratio is insensitive to Te [22]
and does not depend on the elemental abundance. It is,
however, strongly dependent on the electron density N^
(1)
where &EgU is the energy of the transition from g to
u, AQ/H is the oxygen abundance relative to hydrogen,
q(Te) the collisional excitation rate coefficient, Bgu is
the Einstein absorption coefficient, /© the O VI diskaveraged line intensity, D(To,W) accounts for Doppler
dimming and geometrical dilution factors and is function
of the oxygen temperature TO (consisting of the components parallel TO and perpendicular T^~ to the direction
of the magnetic field lines) and the wind velocity W. Ne
is the electron number density, R(Te) is the oxygen ionic
fraction calculated in ionization equilibrium and 0.85 is
the value of the hydrogen to electron number density ratio for a fully ionized plasma with composition given
by [13]. The quantities in brackets (...) are integrated
along the line of sight. We are interested in obtaining
information on the abundance through the comparison
of observed O VI line intensities with values computed
using Eq. 1. Both radiative and collisional components
are strongly dependent on the electron density (N^), electron temperature (Tg), and oxygen elemental abundance
(Aei), while the radiative component depends also on the
adopted values of/ 0 , TO , T^~ and W, We are also dealing
67
10C
10'
fi
O
Q)
~
* 106
1.0
1.5
2.0
Radial Distance in R 0
2.5
3.0
FIGURE 3. Electron density as a function of altitude in polar coronal holes. Measurements obtained in interplume: crosses [14],
open circles [3]. Measurements obtained in plumes: open squares [3]. Measurements obtained without distinguishing between
plume and interplume: open diamonds [15], open triangles [16]. Dashed-dotted line shows the expression, valid from 2 to 4 solar
radii, given by [17]. Dashed and solid lines represent the adopted density profiles in plumes and interplumes, respectively.
3.2
3.0
CO
o
2.8
2.6
C\2
CO
o
2.4
2.2
2.0L
1.4
Radial Distance in R 0
1.6
1.8
FIGURE 4. O VI 1032 to O VI 1037 intensity line ratios as a function of altitude in polar coronal holes. Values obtained in
interplume regions: open circles [present paper], open triangle [18], open squares (UVCS, 21 May 1996) [4]. Values obtained in
plumes: filled circles [present paper]. Solid and dashed lines represent models for interplume and plume conditions, respectively.
Dotted line represents a mixed model comprising plume and interplume lanes.
68
0 VI line intensities in interplume
as well as on the wind speed W. In order to reproduce the
observed line ratio, the Ng profile as a function of height
needs to be known, together with the exciting transition
region radiation. Densities in coronal holes have been
evaluated by several authors [e.g. 17, 23]. However,
in the present analysis only data obtained in 1996 (/. e.
closer in time to our data) were used.
Figure 3 shows Ne as a function of altitude in coronal holes from previous measurements, together with the
adopted density profiles in plumes and interplumes. Using those electron density and temperature profiles the
O VI 1032 and 1037 line intensities were computed using Eq. 1 and the O VI 1032 to 1037 line ratio was then
evaluated. Figure 4 shows that our choice of Ne is able
to reproduce the observed line ratios in interplumes with
negligible outflow velocity up to 1.3 R0. Above this altitude, the comparison with UVCS data obtained on 21
May 1996 in interplume lanes by [4] shows that the region where the fast solar wind is accelerated lies above
1.3 solar radii, as already suggested by [2]. An upper
limit to the outflow speed, possibly present at 1.3 R©, can
be derived by computing the outflow speed which would
bring the O VI 1032/1037 theoretical line intensity ratio
down to the lowest value compatible with the error bar of
the observed ratio. It turns out an outflow velocity of 75
km s"1. We notice, however, that such an high outflow
speed is not consistent with the O VI line ratio observed
by UVCS at 1.5 R0.
Note that W ~ 0, together with the small temperature
gradient, implies the assumption of ionization equilibrium being valid at low altitude above the limb. We are,
hence, confident in using Eq. 1 for calculating line intensities in the low corona.
Using the adopted interplume Te and Ne profiles, we
are able to reproduce the observed intensities in interplumes (Figure 5) with an oxygen abundance of 8.5.
However, this value is lower than that measured in the
fast solar wind, where an abundance of ~ 8.65 ± 0.18
dex is observed in the fast wind around 700 km s"1 ([see
24]). It is interesting to note that the latter value agrees
very well with the recently measured value of 8.69 ± 0.05
dex of the photospheric oxygen abundance ([25]). We
point out that a higher value of Te would translate in a
higher value of abundance, thus allowing us to consider
the value of 8.5 as a lower limit to the value of the oxygen abundance in the interplume lanes of this polar coronal hole. Similar considerations hold also in the case of
the electron density. Higher values of Ne would result in
a theoretical ratio lower than the observed one, while a
lower Ne would result in a small amount of outflow velocity and a higher value of the oxygen abundance.
A realistic attempt to model our plumes observations
requires a combination of plume and interplume emissivities along the line of sight, thus introducing additional
free parameters such as the plume filling factor and oxy-
0 VI 1032
0 VI 1037
10.0
o'.V
1.0
••••••
••_
O
n
O
0.1
1.10
1.40
1.20
1.30
Radial Distance in R Q
FIGURE 5. O VI line intensities as a function of altitude in
an interplume region of the North polar coronal hole. Filled
circles and squares represent the O VI 1032 and O VI 1037
observed line intensities. Open circles and squares represent
the O VI 1032 and O VI 1037 modeled line intensities for the
adopted interplume conditions.
0 VI line intensities in plume
10.0
•
0 VI 1032
•
0 VI 1037
•°*.
D •
ff
1.0
•O
•
O
•n
d
• n
0.1
1.10
1.20
1.30
Radial Distance in R 0
1.40
FIGURE 6. O vi line intensities as a function of altitude in a
plume region of the North polar coronal hole. Filled circles and
squares represent the O vi 1032 and O vi 1037 observed line
intensities. Open circles and squares represent the O vi 1032
and O VI 1037 modeled line intensities obtained by combining
plume and interplume models.
gen abundance. A mixed plume-interplume model has
been, hence, created assuming a single plume embedded
in interplume plasma. The line of sight fraction occupied
by the plume was estimated from Fig. 1 while the same
oxygen abundance as in interplume was adopted. Despite
the fact that we do not feel to attach a physical significance to adopted plume abundance, it is interesting to
note that the introduction of the surrounding interplume
69
plasma is necessary for reproducing the observed line ratios and intensities (see Figures 4 and 6).
6.
7.
CONCLUSIONS
8.
Our measurements of the O VI 1032/1037 line intensity
ratio shows that negligible outflow velocities are consistent with the vaues measured below 1.3 RQ, while an
oxygen abundance of 8.5 is required to reproduce the
observed O vi line intensities in interplume. This value
is lower than the oxygen abundance measured in the
fast solar wind (~ 8.65 ± 0.18). In order to reproduce
the abundances measured in situ, lower N^ values would
be required. Simulations show that densities should be
lower by more than a factor of two than those adopted.
However, such lower Ne profile may imply the solar wind
to be accelerated just above ~ 1.2 solar radii. An indication of this could be found in the apparent decrease of
the observed O vi ratio in interplume above 1.2 R0 but
is apparently inconsistent with the UVCS observations,
which, however, refer to structures observrd about 10
days before the SUMER measurements discussed here.
Because densities play a fundamental role in determining where the solar wind starts being accelerated and in
determining the correct value of the oxygen abundance,
we like to point out that it is of the utmost importance
to have simultaneous determinations of electron density,
temperature and line intensity.
9.
10.
11.
12.
13.
14.
15.
16.
17.
18.
19.
20.
21.
ACKNOWLEDGMENTS
22.
L. T. and G. P. are partially supported by MURST. The
authors would like to thank the anonymous referee for
useful comments and suggestions. Research at Armagh
Observatory is grant-aided by the N. Ireland Dept. of
Culture, Arts and Leisure. The SUMER project is financially supported by DLR, CNES, NASA, and PRODEX.
SUMER is part of SoHO, the Solar and Heliospheric Observatory of ESA and NASA.
23.
24.
25.
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