Solar Wind Composition
Robert F. Wimmer-Schweingruber1
Physikalisches Institut, Universität Bern, Sidlerstrasse 5, CH-3012 Bern, Switzerland
Abstract.
Accurate knowledge of the composition of the solar wind and of solar system abundances allows us to
improve our understanding of various processes affecting the elemental, isotopic, and charge-state abundances
of the solar wind. During the evolution of the Sun, isotopic and elemental abundances have been affected by
migration processes at the interface between the well mixed outer convective zone and the radiative zone. In the
solar atmosphere, the isotopic and elemental composition can be affected by Coulomb drag and wave-particle
interactions. In addition, elemental abundances are clearly fractionated by some process that is controlled by
the first ionization potential (FIP) of the elements. The abundances of the charge states are determined by
processes in the corona, and, in interplanetary space, serve as valuable tracers for the coronal origin of the
solar wind. The abundances of certain key elements and their isotopes can safely and accurately be determined
from their meteoritic values. Their abundances in the solar wind can be used to obtain estimates for the degree
of fractionation undergone by other elements and their isotopes, for which the meteoritic abundances do not
necessarily reflect photospheric values, or for which it is not clear which meteoritic component corresponds to
solar values. In short, the composition of the solar wind can yield valuable information about processes in the early
solar system and its evolution, the evolution of the Sun, processes in the solar interior, atmosphere, and corona,
as well as in interplanetary space. In addition, solar (wind) composition has implications for for astronomy and
astrophysics, for stellar models, and for models of galactic chemical evolution.
INTRODUCTION
does the Sun appear to be chemically equally evolved
as the local interstellar medium [1]? Why does the isotopic composition of non-volatile elements in the galactic cosmic rays agree so well with their solar values [2]?
We would expect the galaxy and hence the interstellar
medium to have evolved chemically since the formation
of the solar system some 4.57 109 years ago.
Nevertheless, the badly understood, but good, agreement of solar and galactic abundances are an improvement to what appeared as an even greater mystery just
about 10 years ago, the “old Sun problem”. Compared
to the local neighborhood, the Sun was metalrich and
hence presumed to be older. Galactic gradients of element abundances are expected to be negative, i.e. the
abundances of heavy elements generally decrease with
increasing distance from the galactic center. Part of the
confusion (but not all) in the past has been due to badly
known solar abundances. While the old Sun problem
has been partly resolved by a finally available volumelimited sample of solar-like stars, recent corrections in
solar metallicity (to which the oxygen abundance is the
prime contributor) also contributed to its resolution [3].
This underlines the importance of determining true abundances (i. e. of some element X relative to H, X/H) as
opposed to the more traditional abundance ratios with
Knowledge of the composition of the Sun and the solar
wind is important for other fields than just solar wind
studies. Apparent changes in solar metallicity with improving measurements have implications for models of
the chemical evolution of the galaxy and for the formation of planet-bearing stellar systems, to name two important examples. Moreover, because the Sun is the star
we can study best of all, it is a unique test bed for theories
of stars and their winds. In addition, detailed studies of
solar system and solar abundances may tell us how to interpret abundances measured in other stars. Can they be
taken straightforwardly, 1:1 from spectra? What corrections need to be made? Studies of solar abundances are
of immense help in interpreting abundances throughout
the galaxy and probably beyond.
Inspite of the fact that we know solar abundances to
a remarkable degree of accuracy (that we will discuss
below), there are several puzzles in solar and galactic
composition that currently elude our understanding. Why
1 now at: Institut für experimentelle und angewandte Physik, ChristianAlbrechts-Universität zu Kiel, D-24098 Kiel, Germany
CP679, Solar Wind Ten: Proceedings of the Tenth International Solar Wind Conference,
edited by M. Velli, R. Bruno, and F. Malara
© 2003 American Institute of Physics 0-7354-0148-9/03/$20.00
577
abundances by nearly 10% [see e. g. 5, for a detailed
comparison with meteoritic abundance ratios]. So just
how representative are photospheric abundances of those
of the bulk Sun and the solar system?
In the early phases of the formation of the solar system, the Sun most likely underwent a violent T-Tauri
phase. Just how did the protoplanetary disk and the
nascent Sun interact? In this phase, a few 10 6 solar
masses of disk material are processed per year. Some material may have been expelled from the solar system by
an X wind [7], some was funneled into the Sun, nonvolatile material may have been ejected back into the
disk along ballistic trajectories, allowing the formation
of chondrules. Given this violent youth, just how representative of solar abundances is the bulk disk material?
Let me illustrate one more question that may be solved
with more precise abundance measurements. Butler et
al. [8] found that planet-bearing stars had a higher metallicity than the Sun, see Figure 1. So far, only stars with
respect to some favorite element, e. g. O or Si. Intriguingly, the old Sun problem still appears to be live and
kicking when one considers isotopic galactic gradients.
There the Sun still lies off the trend expected from galactic chemical evolution models based on measurements of
galactic isotope gradients. Will this again be solved by an
apparent “change” in solar abundances, this time isotopic
abundances?
Whatever the answers to the mentioned questions may
turn out to be, they have certainly motivated why the
study of solar and solar wind composition is important.
We will now proceed to discuss just how well we know
solar and solar wind abundances.
SOLAR SYSTEM ABUNDANCES
As the solar system formed some 4.57 109 years ago, it
closed itself off from the interstellar medium as it existed at that time and at that location. While the bulk
of the matter was concentrated in the Sun itself, only
a small fraction of it is available to us for composition
studies. From characteristic differences in the composition of various phases or separates of different solar system bodies we have been able to acquire a good understanding of the origin of the solar system. For instance,
among the nearly 23’000 cataloged meteorites, there is
one class, the CI chondrites, that exhibit a composition
that is nearly identical to photospheric. Therefore, that
class is considered the most primitive class and often
used to determine solar abundances. However, there are
only five CI chondrites, of which there is an abundant
amount of sample material from only one, Orgeuil [4].
This should illustrate that we should be cautious when
equating solar and solar system abundances. This is often done because meteorites can be studied in the laboratory and hence to a high degree of accuracy. However, it is often forgotten that this accuracy is limited
by sample-to-sample variations of 3 - 10% [H. Palme,
pers. comm.,2002] [see e. g. 5, for an illustrative figure],
while photospheric abundance determinations are limited in accuracy to 25 - 30% by unknowns in the processes on the solar surface and in atomic parameters
[3, 6]. Despite the small number of CI chondrite samples, let us not forget that one gram of such material contains many more atoms than have been analyzed by insitu composition experiments during the space age.
When the accuracy of abundance determinations will
improve sometime in the future (e. g. with the sample
return from the Genesis mission), there will be a new
set of problems that will need to be investigated. For
instance, the interplay of radiative forces, collisions, and
gravitation in the (time-evolving) outer convective zone
of the Sun leads to a depletion of the heavy element
10
Sun
field stars
planet−bearing
stars
Number of stars
8
6
4
2
0
−0.6
−0.4
−0.2
0.0
0.2
Metallicity [Fe/H]
0.4
0.6
FIGURE 1. Comparison of the metallicity of planet-bearing
and field stars. Data are from Butler et al. [8]. Plotted are the
number of field stars (shaded, thin line) and planet-bearing stars
(empty, thick line) in a given metallicity bin. Metallicities are
normalized to solar and given on a logarithmic scale (dex).
at least Saturn-sized companions have been detected [8].
Smaller planets cannot presently be detected because
their influence on the central star is too small. Because
giant gas planets only form beyond the ice line, their
detection closer to their star means that they have migrated inward. And any Earth-like, rocky planet must
have migrated too far. . . Planet-bearing stars also show
a striking surface enhancement of Li [9], an element that
is largely destroyed by nucleosynthetic processes in the
early phases of the evolution of Sun-like stars. That Li is
depleted in the solar photosphere would then imply that
the Sun did not “eat up” a substantial amount of planetary
matter after the formation of an outer convective zone.
This implies that the disk cleared faster than the migra-
578
tion time for terrestrial planets, possibly due to intense
radiation of neighboring giant stars which were born in
the same star-forming region as the Sun and which may
have contributed to or altered the composition of disk
material.
of other ionizing scenarios, e. g. electron impact ionization, have not been computed to my knowledge, and we
do not know how they would organize a FIT plot.
Other fractionation mechanisms such as inefficient
Coulomb drag or wave-particle interactions further modify the abundances of the nascent solar wind from chromospheric to what we measure in interplanetary space
[17, 18]. A good understanding of these effects is necessary for an accurate determination of solar abundances
from solar wind measurements. With the sample-return
mission Genesis now collecting solar wind (see the paper
by Reisenfeld et al. presented at this conference) we need
to address fractionation mechanisms to a much higher
degree of accuracy [See e. g. 19, for a review.].
SOLAR WIND COMPOSITION
The composition of the solar wind is one step removed
from that of the photosphere. It is altered by a multitude
of fractionation mechanisms which may not all be active
at the same time and locations. The most striking effect
is the fractionation according to first ionization potential
(FIP) which appears to enhance the abundance of elements with low FIP (below about 10 eV) with respect
to elements with a higher FIP. However, the previously
clear cut FIP step between low- and high-FIP elements
does not appear to be as clear as it used to be. Plotting
the familiar FIP plot with logarithmic axes results in the
plot shown in Figure 2. A dashed line (∝ FIP 1 ) has been
drawn to guide the eye and appears to organize the plot
just as well as a step-like curve.
Al
(X/O)SW/(X/O)photo
So far, we have considered the use of solar (wind) composition for what might be termed “studies of origins” the origins of the Sun and its planets, the solar system,
or the evolution of the galaxy. Of course, certain models
need to be developed to make corrections for fractionation mechanisms or other processes that could modify
the abundances of the elements. One of the beauties of
composition studies is the possibility to use composition
measurements to test these models, to infer the importance of the various processes, but also to to use the many
possible combinations of elements, isotopes, and ions for
studies of the underlying plasma physics which determines the properties of the solar atmosphere, the heliosphere, and beyond. Let us consider just a few examples
of how changes or differences in solar wind composition
can be used to infer the importance of various processes
active in the chromosphere, the corona, and interplanetary space.
Xe
Fe
Mg Si
slow
fast
soil
Ca
Na
SOLAR WIND COMPOSITION AS A
TRACER
S
Kr
Ar
H
1
C
Ne
O
He
N
10
First Ionization Potential (FIP) [eV]
FIGURE 2. Comparison of solar wind abundance ratios normalized with O with their photospheric values. Elements with
low first ionization potential (FIP) are enhanced with respect to
high-FIP elements. The dashed horizontal line indicates photospheric values. The diagonal dashed line corresponds to an
abundance enhancement ordered by 1/FIP and forced to be
unity for O. It is only drawn to guide the eye. Solar wind abundances were taken from the review of [10] and normalized to
the photospheric abundances of [11]
The Chromosphere
The chromosphere appears to be the most likely site
of the FIP effect. Relating abundance enhancements in
a FIP plot to FIT allows us to determine timescales for
the atom-ion separation process, typically a few to a
few 10s of seconds. The process gradually depletes the
high-FIP elements until the material is released to the
solar wind. Widing and Feldman [20] have investigated
the FIP enhancement in active regions over time. They
found that the material in the loops is unfractionated
right after appearance of the loop and that the ratio of
Mg/Ne (their proxy for low FIP to high FIP elements)
increases approximately linearly with time. Hence solar
The origin of the FIP effect seems to lie in some mechanism that separates atoms and ions [12, 13, 14]. Measurements of noble gas abundances in lunar soils [15]
may help to better understand this mechanism because
they do not appear to follow the same trend as the other
elements in Figure 2. Assuming solar matter to be ionized by UV and EUV photons, one may derive a first ionization time (FIT) which appears to order all elements,
including Ar, Kr, and Xe, better than FIP [16]. The FITs
579
different loop heights [28] and is experimentally found
to result from physically sensible parameters [29]. It is
not clear, however, how the systematic ordering of and
large differences in freeze-in temperatures arises in this
scenario. If they are established higher up in the corona,
as is generally assumed, then how should the plasma retain its memory of the loop temperature or loop height?
If the temperature is set in the loops, why do charge-state
ratios of C, O, Mg, Si, S, Fe, etc. all show different temperatures?
wind abundance measurements should give information
about the lifetime of loops prior to disruption.
The Corona
On its way from the photosphere to interplanetary
space, the solar wind is further modified in the corona
e. g. by collisions with coronal electrons. The chargestate composition of an element is determined by by
a competition of timescales. As the solar wind is accelerated through the corona it experiences a diminishing electron number density and hence an increase in
the ionization timescale, but a decrease in the expansion timescale because the solar wind is being accelerated. Where the two timescales are equal, we traditionally say that the charge states freeze in [21] at that location and electron temperature. Unfortunately, this picture
is overly simplistic and one cannot determine a coronal
electron temperature profile from charge-state measurements in any straightforward way (See the paper by Esser
et al., presented at this conference.). Not all ions experience the same acceleration and hence differential streaming between different ions tries to establish itself. This is
counteracted by the change of identity of ions when their
charge state is modified by ionization or recombination.
Depending on where in the corona the ions form, differential streaming is an important factor in determining the
charge-state composition of the solar wind [22].
Interplanetary Science
In-situ measurements of solar wind heavy ions can
give information about kinetic processes in the heliosphere. Early measurements of the speeds and kinetic
temperatures of protons and alpha particles showed that
the latter often flow faster then the former and that
the temperatures are mostly mass proportional [see e. g.
30, 31, for a review]. Early measurements of heavy ions
showed that they generally flow at a speed in between
the proton and alpha-particle speed [32, 33, 34, 35, 36],
a finding that has been confirmed recently [37]. The
mass proportionality of kinetic temperatures also holds
for heavy ions at high temperatures, however, the cold
solar wind appears to be isothermal [37]. These results
must be compared with the observation of von Steiger
et al., who found that, at 5 AU heliocentric distance,
all heavy ions have kinetic temperatures which are mass
proportional. In other words, the transition from collision
dominated (isothermal) cold solar wind to wave heated
(mass proportional temperatures) solar wind takes place
between 1 and 5 AU. At 1 AU slightly warmer wind
lies in a regime between collision and wave-particleinteraction dominated. In other words, 1 AU is an interesting place for in-situ instruments because that is where
this transition takes place. Because wave-particle interaction is sensitive to the mass and charge of the interacting ions, and heavy ions contribute a large number of
mass numbers and charge states, accurate measurement
of their 3-d velocity distribution functions in various solar wind regimes for instance by the PLASTIC instrument on STEREO will be very interesting.
Recent Developments
The recently proposed model for the heliospheric
magnetic field around solar activity minimum [23] not
only is a possible explanation for the transport of energetic particles to high heliographic latitudes, but also
for the quasi-rigid rotation of coronal holes and possibly even for the origin of the solar wind. This model
differs from similar earlier work of Wang and Sheeley
[24] by the introduction of a systematic and self consistent transport of field lines across the corona. An interesting consequence of this transport is the necessity
of reconnection and migration of open field lines in the
coronal hole and in the streamer belt. The reconnecting
field lines open up previously closed loops, setting free
the enclosed plasma, thus giving rise to the solar wind. In
this picture the older and larger loops in the streamer belt
have had more time to fractionate (see subsection chromosphere) and, by virtue of being higher than in coronal
holes, should exhibit higher electron temperatures than
the smaller and younger loops in coronal holes. The well
known anti correlation of freeze-in temperatures with solar wind speed [25, 26, 27]is then interpreted as due to
Coronal Mass Ejections
Coronal mass ejections (CMEs) are often but not always associated with compositional peculiarities. Large
abundances of alpha particles and sometimes of singly
charged helium have been reported [38], sometimes associated with high charge states of iron [38], even Fe20
has been reported [39]. Occasionally, large overabun-
580
dances of the very rare isotope 3 He are observed [40],
some spectacular CMEs have very unusual composition,
such as the May 1998 event [41, 42, 43]. A class of
CMEs appears to exhibit evidence for mass-proportional
fractionation, possibly due to leakage of material out of
loop structures before the eruption [44, 45, and the article by Wurz et al. in these proceedings.]. In a study of
42 CMEs observed with SWICS on Ulysses, Neukomm
[46] found that CMEs can be divided into three classes,
one similar to the slow and the fast wind, respectively,
and one distinct from any type of solar wind. Such surveys are important, because they drive home the point
that not all CMEs are unusual, in fact, typical CMEs are
quite unspectacular composition-wise. We often tend to
concentrate on some “interesting cases” because they are
spectacular - and then to forget that they are spectacular, not typical. Zurbuchen et al. give a summary of CME
composition in these proceedings.
The composition of CMEs is probably determined in
two phases with quite different time scales. The elemental (and isotopic) composition is largely determined
in the time preceding the eruption. During this phase,
various processes can fractionate elements according to
atomic properties such as FIP, mass, or mass per charge.
Recalling the results of Widing and Feldman [20], we
should be able to determine how long the CME material
was contained in an active region loop prior to eruption
(it would also be interesting to compare with the Wurz
model [45]). After the CME has begun to rise through the
solar atmosphere, the charge states of the many different
elements are affected by the rapidly changing electron
energy distribution function.
such as co-rotating interaction regions (CIRs) [47, 48].
The structure of CMEs is also (at least partially) resolved in composition measurements [e.g. 44, 42]. High
time resolution measurements of velocity distributions of
heavy ions will give new information about plasma processes active in the inner heliosphere.
Accurate knowledge of solar abundances are important for a variety of applications, not just in our field, but
also in astronomy and astrophysics. Most important, the
Sun is the only star which we can study in the detail to
which we have studied the Sun. For instance, here we can
test model predictions for abundance alterations by comparing solar system with solar abundances. This cannot
be done for other stars. Thus the solar and heliospheric
community is in a unique position and communication
between the solar and stellar composition communities
needs to be increased.
ACKNOWLEDGMENTS
It is a pleasure to acknowledge stimulating discussions
with P. Bochsler, P. Wurz, R. v. Steiger, H. Holweger,
G. Gloeckler, T. H. Zurbuchen, and the many participants
of the Joint SOHO/ACE workshop on “Solar and Galactic Composition”. Furthermore, I wish to thank the organizers of SW 10 for their flawless organization of a very
inspiring conference. This work was supported in parts
by the Swiss National Science Foundation and the Canton of Bern.
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SUMMARY, DISCUSSION, AND
CONCLUSIONS
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