Show Summary Details

Page of

Printed from Oxford Research Encyclopedias, Planetary Science. Under the terms of the licence agreement, an individual user may print out a single article for personal use (for details see Privacy Policy and Legal Notice).

date: 15 January 2021

Solar Elemental Abundancesfree

  • Katharina LoddersKatharina LoddersEarth and Planetary Sciences and McDonnell Center for the Space Sciences, Washington University in Saint Louis

Summary

Solar elemental abundances, or solar system elemental abundances, refer to the complement of chemical elements in the entire Solar System. The Sun contains more than 99% of the mass in the solar system and therefore the composition of the Sun is a good proxy for the composition of the overall solar system. The solar system composition can be taken as the overall composition of the molecular cloud within the interstellar medium from which the solar system formed 4.567 billion years ago. Active research areas in astronomy and cosmochemistry model collapse of a molecular cloud of solar composition into a star with a planetary system and the physical and chemical fractionation of the elements during planetary formation and differentiation. The solar system composition is the initial composition from which all solar system objects (the Sun, terrestrial planets, gas giant planets, planetary satellites and moons, asteroids, Kuiper-belt objects, and comets) were derived.

Other dwarf stars (with hydrostatic hydrogen-burning in their cores) like the Sun (type G2V dwarf star) within the solar neighborhood have compositions similar to the Sun and the solar system composition. In general, differential comparisons of stellar compositions provide insights about stellar evolution as functions of stellar mass and age and ongoing nucleosynthesis but also about galactic chemical evolution when elemental compositions of stellar populations across the Milky Way Galaxy is considered. Comparisons to solar composition can reveal element destruction (e.g., Li) in the Sun and in other dwarf stars. The comparisons also show element production of, for example, C, N, O, and the heavy elements made by the s-process in low to intermediate mass stars (3–7 solar masses) after these evolved from their dwarf-star stage into red giant stars (where hydrogen and helium burning can occur in shells around their cores). The solar system abundances are and have been a critical test composition for nucleosynthesis models and models of galactic chemical evolution, which aim ultimately to track the production of the elements heavier than hydrogen and helium in the generation of stars that came forth after the Big Bang 13.4 billion years ago.

Definitions

Element concentrations usually refer to the elemental content of a given system by mass fraction or mass percent of total system mass present. Common units are gram per gram (g/g) units, where 0.01g/g=1weight%=1mass%, milligram per gram (mg/g=103g/g=0.1mass%), microgram per gram (μ‎g/g) or parts per million (ppm) (1μg/g=106g/g=1ppm=0.0001mass%), nanogram per gram (ng/g) or parts per billion (ppb), where 1ng/g=109g/g=1ppb=0.001ppm, and parts per trillion (ppt) (1012g/g = 0.001 ppb).Sometimes atomic or molar concentrations are used; these refer to the number of atoms or molecules as part of total atoms or molecules present. Common units are atom percent (at%) or mole fractions.

Elemental abundances refer to the relative elemental contents of a given system (e.g., the Sun, a planet, meteorite, or comet). Abundances are typically reported on an atomic scale by number relative to the number of hydrogen atoms (astronomical abundance scale) or silicon atoms (cosmochemical abundance scale). Relative abundances by mass are rarely used. Sometimes the terms abundances and concentrations are used interchangeably. Concentrations refer to content as a part of the total system, and abundances refer to content in relation to contents of a reference species.

Regarding solar mass-fractions X, Y, and Z of the elements, the mass fraction of hydrogen and helium in the Sun are called “X” and “Y,” respectively. The symbol “Z” is the mass fraction of all other elements heavier than helium (Li to U). Astronomers lump together all heavier elements because these only make up about Z=0.0148 of the Sun whereas hydrogen has X=0.7389 and helium Y=0.2463. The solar mass fractions are important quantities for the standard solar model that describes how the Sun works and evolved over time. The mass fractions quoted here are for the outer convection zone of the Sun and are based on helioseismic measurements (Basu & Antia, 2004), and the Z/X ratio (0.020) is derived independently from the abundances in the photosphere (see later discussion), which is the observable top layer of the Sun’s convection zone.

Solar system abundances refer to the elemental inventory of the solar system at the time of solar system formation 4.567 billion years ago. This includes the Sun, the planets and their moons, asteroids, comets, meteorites, and interplanetary dust. Historically, the abundances are mainly derived from meteorites and spectroscopy of the solar photosphere. Sometimes solar system abundances are referred to as “solar abundances,” but technically there are important differences in how solar and solar system abundances are derived.

Solar or solar photospheric abundances refer to the elemental inventory of the spectroscopically accessible parts of the outer convective envelope of the present-day sun, where “present day” is in respect to the young Sun 4.6 billion years ago. In addition to the photosphere, solar abundances are obtained from sunspots, the solar corona, and the solar wind. In the photosphere the elements are mainly present as neutral or singly ionized atoms; a few absorption bands from C2, CH, CN, NH, CO, and OH are known. These molecules occur in larger quantities in sunspots in the cooler sunspots where MgH, CaH, TiO, BO, HS, HF, and HCl are also present. Photospheric abundances reflect the composition of the Sun’s outer convection zone (about 2% of the Sun’s total mass) and are not representative of the Sun’s composition 4.6 Ga ago because of diffusion and gravitational settling into the Sun’s interior over the Sun’s lifetime. Roughly 10% to 20% (depending on models adopted; see later discussion) of each element heavier than hydrogen was lost from the photosphere while hydrogen diffused up. The Sun is losing mass (about 1014 solar masses per year) through the Sun’s outer layers in the form of solar wind. Assuming about similar mass loss in the past, this mass loss from the convection zone (about 0.2% total of the current convection zone mass) is negligible for the Sun. The elemental abundances in the solar wind are fractionated relative to photospheric values, and mass loss from the Sun must contribute to relative element abundance fractionations within the photosphere over time, although on a much smaller scale than fractionations stemming from element settling.

Meteoritic abundances are derived from the composition of meteorites. In the context of solar system abundances, meteoritic abundances typically refer to the abundances of condensable elements derived from CI-chondrites (see later discussion and the article on meteorites in this encyclopedia). In the older literature, meteoritic abundances were often derived from mass-weighted average compositions of the silicate, metal, and sulfide phases present in chondritic meteorites (for chondrite classification, see Table 1). This approach suffered from the necessary assumption of representative phase compositions and the values for representative mass fractions of silicate, metal, and sulfide corresponding to the initial solar system composition.

Abundance Scales

The astronomical abundance scale uses hydrogen, the most abundant element in the Sun, as a reference element. Relative abundances in the photosphere or other stars are frequently given in the epsilon notation:

log10εE=log10(E/H)=log10(n(E)/n(H))(1)

for the atomic ratio of an element “E” relative to hydrogen (H). Sometimes the element symbol is taken as shorthand to express the number of atoms, n(E). Another frequently used notation is

A(E)=12+log10εE=12+log10(n(E)/n(H))(2)

which expresses abundances on a logarithmic scale relative to n(H)=1012 atoms. The predominance of hydrogen in the Sun and other stars was recognized in the 1920s, and Cecilia Payne (1925a, 1925b) found A(H)>11 for the Sun and A(H)=12.5 for her set of stars analyzed. In her thesis, Payne (1925b) defined the fractional log10 abundance to A(H)=11. The normalization with hydrogen to A(H)=12 was apparently first done by Claas (1951) in order to avoid negative logarithmic values for elements of low abundance. Frequently the term “dex” is used in the astronomical literature to designate “decadic logarithmic units” where 1 dex stands for a factor of 10(=101).

Atomic abundances on the cosmochemical scale are frequently normalized to the number of silicon atoms set to N(Si)=106. Abundances on this scale are sometimes designated as N(M). These abundances can be calculated from the absolute concentrations by mass.

Goldschmidt (1937) introduced the normalization to silicon, the most abundant cation in the crust (oxygen is typically the most abundant element by number in the crust and rocks but is difficult to measure routinely). Cations are a better choice to normalize the rock-based abundance scale. Originally, this scale was set to N(Si)=100 (Goldschmidt, 1938), then N(Si)=10,000 (Brown, 1949). With more refined analyses, smaller quantities of many elements could be measured, and it became more practical to increase the normalization. In their seminal paper, Suess and Urey (1956, p. 57) changed the scale to N(Si)=106, explaining, “We use 106 in order to get values for the rarer elements which can be written without negative exponentials or awkward decimal fractions.”

For comparing the astronomical and cosmochemical abundance scales, renormalizing meteoritic values to H is inappropriate because H forms volatile compounds and is highly depleted in meteorites. It is more practical to normalize photospheric abundances (which are always relative and not absolute) to Si=106 atoms. Any other well-determined refractory elements can serve to link the meteoritic and photospheric scales, and other elements such as Ca (Holweger, 2001), Mg (Goles, 1969; Larimer & Anders, 1967), and Ti (Greenland & Lovering, 1965) have been suggested over time.

The abundances on the astronomical log scale, A(E) are related to the cosmochemical scale, n(E) using the relationships:

A(E)=AF+log n(E)(3)

or

n(E)=10A(E)/10AF=10(A(E)AF)(4)

The term “AF” is often called “abundance scaling factor” and determined by comparing elements that are well determined in the Sun and in meteorites.

Sometimes only one element (traditionally silicon) is used to determine the scaling factor (e.g., Lodders, 2003). On the other hand, this requires recomputation of all elemental abundances if the abundance for silicon for meteorites and/or in the photosphere is changed. For example, Lodders, Palme, and Gail (2009) recommended a photospheric Si abundance for which Asplund, Grevesse, Sauval, and Scott (2009) found a value different by 0.02 dex, and Asplund et al. (2009) then changed all meteoritic abundances from Lodders et al. by 0.02 dex to obtain their meteoritic values on the astronomical log scale. On the other hand, a mean scaling factor derived from several elements typically makes the coupling of the scales more robust and less sensitive to changes in individual elemental abundances. This approach was introduced by Urey (1967) and Cameron (1968), who used 15 and 10 nonvolatile elements, respectively. This practice has been widely followed (e.g., Anders & Grevesse, 1989; Lodders et al., 2009; Palme & Beer, 1993; Palme, Lodders, & Jones, 2014). Typically, the scaling factor derived from several elements has an uncertainty of up to ±0.04 dex (an absolute factor of 1.1), which should be considered when gauging how well meteoritic and solar abundances agree once placed onto the same relative abundance scale. Using a scaling factor from several elements may also lead to abundances scales that are nominally normalized to Si=106 but where the actual silicon abundance does not equal 106 (e.g., in Anders & Grevesse, 1989). Here the recommended values use the scaling factor from silicon to obtain clean, normalized scales. The photospheric and meteoritic abundances on the two atomic scales are summarized in Tables 4 and 6, and recommended solar system abundances on both scales are given in Table 8.

The uncertainties of the logarithmic abundance scale and linear abundance scale are compared using the relationship U(%)=±100(10±a1) where “a” is the uncertainty in decadic logarithmic “dex” units quoted for photospheric abundances on the logarithmic scale, and “U” is the uncertainty of the meteoritic abundance in percent. The uncertainty in dex is an uncertainty factor, hence the percent uncertainty is smaller for –a than for +a. Vice versa, a given percent uncertainty on a linear scale yields two different uncertainty factors on the logarithmic scale. For a conservative approach, the larger percentage value U (from a given uncertainty “a”) or the larger uncertainty “a” (from a given U) is taken for comparison here. Uncertainties of meteoritic as well as photospheric and stellar abundances on the logarithmic scale are frequently quoted as 1- or 2-sigma confidence intervals; however, the original papers do not always specify which confidence intervals were used for the reported uncertainties. Both statistical and systematic errors must be considered. The statistical deviation in solar abundances from the line-to-line scatter (when several lines are used to determine elemental abundances) is typically much smaller than the uncertainties introduced by parameters of the solar atmospheric model used for the analyses (e.g., Lawler, Feigenson, Sneden, Cowan, & Nave, 2018; Lawler et al., 2014; Scott, Asplund, Grevesse, Bergemann, & Sauval, 2015; Scott, Grevesse, et al., 2015; Sneden et al., 2009; and the review by Allende-Prieto, 2016).

Historical Background

About a century ago, the era of modern quantitative studies about the abundances and distribution of the chemical elements started as almost all naturally occurring elements of the periodic table had been discovered. Chemical analyses were no longer limited to major elements occurring in rocks, and quantum mechanics had opened the door to quantitative spectral analyses of the Sun and stars. There was and remains a strong desire to know which elements exist and their quantities on the surface of the Earth and within the Earth and planets, the Sun, and the cosmos. The nucleosynthetic, geochemical, and cosmochemical principles governing formation and distribution of the elements remain active research areas in astrophysics and cosmochemistry. Traditionally, the knowledge about elemental abundances comes from two disciplines—astronomy with studies about the composition of the Sun and other stars and cosmochemistry with hands-on studies of meteorites, interplanetary, and interstellar dust. Both disciplines complement each other to validate, supplement, and optimize solar system abundance data, which are widely used as a reference composition in both disciplines.

Like others (Chladni, 1819; Farrington, 1915; Harkins, 1917; Merrill, 1909, 1915, 1916) before him, Victor Moritz Goldschmidt (1923, 1937, 1938) realized that stony meteorites seem to be the least differentiated rocky objects compared to the Earth. The Earth has differentiated into a crust, a silicate mantle, and a metallic core and has ongoing geological activity (volcanism, plate tectonics). Goldschmidt and others thought meteorites presumably have been idle on their parent asteroids since shortly after their formation 4.567 Ga ago, yet even the so-called primitive meteorites were exposed to space weathering and often bear witness to impact processes they experienced on their parent asteroids (see, e.g., article in this encyclopedia on meteorites by Wang & Korotev).

Goldschmidt’s work was quite influential as he developed the geochemical classification of the elements in a series of papers on “Geochemische Verteilungsgesetze der Elemente” (Geochemical Distribution Laws of the Elements), which describes the affinities of the elements for certain phases: Rock-loving (lithophile) elements such as Ca, Al, Mg, Si, and, to some extent, Fe usually occur in oxide form in rocks. Other elements such as noble metals, Ni, Co, and Fe appear in metal alloys (siderophile or metal-loving elements), and other elements are preferentially associated with sulfide phases (chalcophile elements such as Zn, S, Se, and Fe). The group of atmophile elements contains the noble gases and elements in compounds present in the atmosphere (C, N, O, H) and are easily volatilized under typical conditions near the surface of the Earth.

Many analyses of meteorites by Goldschmidt and coworkers, his direct competitors Noddack and Noddack (1930), and the solar abundances of H, C, N, and O from Russell (1929) provided Goldschmidt (1937) with the data needed to assemble a comprehensive table of the elemental abundances in cosmic matter. Although Goldschmidt’s table was instrumental for further work on nucleosynthesis, others predated his work and laid the foundations.

In 1914, Russell made four lists of the elements in order of line intensities in the solar photosperic and coronal spectra and of elemental abundances in the Earth’s crust and stony meteorites. He found qualitative resemblances in the ordering for 15 out of the 16 abundant metallic elements between the Sun, crust, and meteorites. In 1914, the large abundances of hydrogen and helium in the Sun and stars were not yet recognized, and Russell (1914, p. 794) noted that, with some exceptions, “the agreement of the solar and terrestrial lists is such as to confirm very strongly Rowland’s opinion that, if the earth’s crust should be raised to the temperature of the Sun’s atmosphere, it would give a very similar absorption spectrum. A moderate admixture of meteoritic material would make the similarity even closer.”

In this context, it is worth adding that de Beaumont in 1847 (before spectral analyses was invented) had already found 16 abundant elements common to crustal rocks, ore deposits, volcanic emanations, mineral and ocean waters, meteorites, and organic matter. Thus, by 1914, Rowland and Russell extended the qualitative comparison to the Sun. Plaskett (1922) and in particular Payne (1925a, 1925b) made this comparison more quantitatively. In her thesis, Payne also recognized the large abundances of H and He (but that result was discouraged by her advisor, Russell). Thus she limited the comparison to abundant elements in common and describes (Payne, 1925b, p. 187) “The relative abundances, in the stellar atmosphere and the earth, of the elements that are known to occur in both, display a striking numerical parallelism.” She further noted the better agreement when stellar compositions are compared to the total Earth composition instead to the Earth’s crust.

By 1929, Russell had accepted the overwhelmingly large abundances of hydrogen and He and presented the first extended and more or less quantitative abundance determinations of 56 elements in the solar photosphere. Russell compared his solar abundance table to meteorite abundances by Merrill (1915) and to abundances in the Earth’s crust (Clarke & Washington, 1922); these abundances also had improved since 1914. Russell noted stronger similarities in the relative abundances of certain elements in the solar photosphere and meteorites than between photospheric abundances and those in the Earth’s crust.

Subsequent abundance tables of the elements became more complete and refined through better analytical techniques for meteorites and solar spectroscopy but also through better theoretical understandings of the underlying principles of atomic and nuclear physics and nucleosynthesis.

Meteoritic Abundances

Many meteorites contain silicates, metal, and sulfides in proportions that one might expect for the Earth and other terrestrial planets, where silicate rocks and metal have separated to form a silicate mantle and crust and a core. As early as Chladni (1794), this observation triggered the idea that meteorites may represent “leftover” material from the birth of the solar system and thus provide information about the elemental composition of the matter from which the solar system objects formed. A brief summary about chondritic meteorites follows, and the reader can find more details about meteorites, including definitions and explanations about their classifications, in the “Meteorite” article of this encyclopedia, in Scott and Krot (2014), Lodders and Fegley (1998, 2011), in the Meteoritical Bulletin from the Meteoritical Society, on the excellent web page about meteorites by Randy Korotev, and in the article by Wang and Korotev in this encyclopedia.

Table 1. Major Chondrite Meteorite Classes

Designation

Classification Notes

Carbonaceous Chondrites

CI

I for “Ivuna” type, old: Carbonaceous I (Wiik, 1956)

CM

M for “Murray” type, old: Carbonaceous II (Wiik, 1956))

CV

V for “Vigarano” type

CO

O for “Ornans” type, old: Carbonaceous III, Ornanites (Wiik, 1956)

CK

K for “Karoonda” type

CR

R for “Renazzo” type

CB

B for “Bencubbin” type

CH

H for “High metal content”

Ordinary Chondrites

H

H for “High iron content” (Urey & Craig, 1953), old, bronzite-olivine, or Cronstad-type (Prior, 1916)

L

L for “Low total iron content” (Urey & Craig, 1953), old, hypersthene-olivine or Soko-Banja type (Prior, 1916)

LL

LL for “Low total iron and low total metal content” (Keil & Fredriksson, 1964; Urey & Craig, 1953), old, hypersthene-olivine or Baroti-type (Prior, 1916), amphoterites (Mason & Wiik, 1964)

K

K for “Kakangari”

R

R for “Rumuruti”

Enstatite Chondrites

EH

EH for “Enstatite Chondrite with High-Fe metal content,” old, HH (Wiik, 1956) and old type I (Anders, 1964)

EL

EL for “Enstatite chondrite with Low-Fe metal content,” old, type II (Anders, 1964)

Rose (1863–1864) recognized the major meteorite chondrite groups and introduced the name “chondrites” for stony meteorites containing the silicate spheres. Chondrites are chemically “primitive” or “undifferentiated”; thatis, they did not melt and partition into silicates, metal, and sulfide phases by density separation. Prior (1916, 1920) established the basics of the classification scheme (Table 1), which was expanded upon by Urey and Craig (1953), Wiik (1956), and Mason (1962) as chemical analyses revealed characteristic differences (see Mason, 1971). The broad classes, based on composition and silicate mineralogy, are carbonaceous chondrites, ordinary or common chondrites, and enstatite chondrites. The carbonaceous meteorite class is subdivided depending on overall carbon and water contents and mineralogy. Wiik (1956) used types “I, II, and III,” but these are now named after a type-specimen meteorite (e.g., the CI-chondrites are named after the meteorite fall near Ivuna, Tanzania). A similar name-based designation existed once for other chondrite classes but was abandoned in favor of a designation from their silicate mineralogy as enstatite chondrites, and for ordinary chondrites as bronzite and hypersthene chondrities (and later additionally amphoterites). The ordinary chondrites, based on their relative metal contents and iron concentrations in the silicates, are H, L, and LL (Keil & Fredrrikson, 1964; Urey & Craig, 1953) and EH and EL for enstatite chondrites. This practice has been used to name a more recently designated group of metal-rich carbonaceous chondrites as “CH,” but another recent class (“CB” for “Bencubbinites”) was named after the Bencubbin meteorite.

The numerical subclassification for the petrologic type by van Schmus and Wood (1967) describes successive mineral alteration that meteorites experienced on their parent asteroid. The most primitive types were designated “type 3” (for which subdivisions from 3.0 to 3.9 were introduced later). Increasing numbers from 4 to 7 describe mineralogical alteration effects from mild to more severe heating close to melting. Numbers below 3 designate mineralogical alteration caused by aqueous fluids, with type 2 having partial mineral transformation and type 1 more or less complete mineral transformation from “dry” to hydrous minerals, and essentially oxidation of metal to magnetite. The review by Scott and Krot (2014) describes chondrites in more detail. None of the meteorite groups has escaped some processing on their parent asteroid, so the chemical and mineralogical constitution of the original matter in the solar system has been altered or lost. However, the elemental composition was retained in cases where isochemical reactions took place, such as generally assumed for CI-chondrites, which were established as the preferred solar system abundance standards for nonvolatile elements in the 1970s. It took some time to investigate and characterize the meteorite classes, and it is no coincidence that the selection of the CI-chondrites culminated at about same time that analytical advances provided the information to solidify the chondrite meteorite classification.

Early studies derived meteoritic elemental abundances from weighted averages of elemental concentrations in meteoritic silicates, metal, and sulfide, irrespective of meteorite type, which meant occasionally counting stony meteorites as “silicate” and iron meteorites as “metal” phase. In part this choice was dictated by the way wet-chemical analyses were done, and the analyses were mainly restricted to more abundant elements. Minor elements could be analyzed only when large quantities of meteoritic metal, silicates, or sulfide were available. Given the scarcity of meteorites, small amounts of samples were subjected to destructive analyses, and only developments in microanalytical techniques brought improvement in quantitative analyses to allow the use of whole-rock (bulk) analyses for deriving solar system abundances from certain meteorite groups. Microscopy for petrological studies, and spectral absorption and emission spectroscopy, and colorimetric analyses had revolutionized chemical analyses since the 1860s, and, similarly, analytical developments now can provide measurements of low atomic concentrations on micro- to nano-scale samples and in some instances even beyond. Over time, analytical tools for elemental and isotopic measurements applied to abundance studies in meteorites include X-ray diffraction, X-ray spectrography, and X-ray fluorescence, mass spectroscopy since the early 1900s, radiochemical methods with neutron-activation analysis, isotope dilution and electron microprobe analyses since the 1950s, secondary ion probe mass spectroscopy since the 1970s, inductively coupled plasma mass spectroscopy (ICP-MS) since the 1990s, and multicollector ICP-MS since the 2000s. Currently much more sensitive analytical techniques (e.g., ICP-MS) boost detection limits toward part per trillion concentrations and provide elemental and isotopic abundance determinations of meteorites and their components.

Having the analytical data leads to the question of how to derive representative abundances from them. Many early estimates of representative meteorite compositions listed in Table 2 used measurements that provided concentrations in the silicate, metal, and sulfide portions. Then, for a representative weighted average, the proportions of silicate, metal, and sulfide had to be assumed. The leading thought was of an original “planet” from which all meteorites originated, and this planet must have had a density much like the Earth and the other terrestrial planets. The approach was to select meteoritic silicate, metal, and sulfide in proportions to match the densities of the Earth or the average terrestrial planet densities. Then the overall meteoritic composition could be calculated if the elemental compositions of each phase were known. Farrington (1915), Noddack and Noddack (1930, 1935), Saslawsky (1933), Goldschmidt (1937, 1938), Brown (1949), Levin, Kozlovskaia, and Starkova (1956), Urey (1952), and Mason (1962) and others followed this procedure. More details about these approaches are in the footnote to Table 2.

The alternative was to assume chondrites with their silicates, metal, and sulfides are already a representative mixture, because their chemical similarities, especially when compared on a volatile-compound free (i.e., no carbon, water, or nitrogen bearing compounds) basis, were noted as early as Wahl (1910). According to Mason (1960, p. 1), chondrites are not leftover debris from a disrupted planet but “have always been independent and individual objects.” The approach of taking grand averages of available data for chondrites as representative was used by Merrill (1909, 1915, 1916), Harkins (1917), and Saslawsky (1933). Noddack and Noddack (1934) used a more pragmatic approach to obtain an average chondrite composition by analyzing a physical mixture of many different pulverized chondrites (but they still adhered to using weighted phase average to derive representative meteoritic compositions). Urey and Craig (1953) adopted chondrites to search out what needed to be done with more modern analysis tools at the time. As more meteorites were analyzed chemically and mineralogically, the question arose regarding which meteorite group is closest to the average composition of condensable elements of the solar system.

Table 2. Early Estimates of Representative Meteoritic Compositions

Z

Average Phases

Average Phases

Average Phases

Average Phases

Stony Mets.

Average Phases

Stony Mets.

Average Phases

Average Phases

Stony Mets.

Stony Mets.

Average Phases

Stony Mets.

[Mas62]

[Lev56]

[U52]

[Bro49]

[Budd46]

[Gold37]

[NN34]

[Sas33]

[NN30]

[Har17]

[Mer16]

[Far15]

[Mer09]

1

H

380

830

220

2

He

3

Li

2

3.2

5

2

5

4

5

3

4

Be

[1]

0.09

1

0.6

10

1

10

5

5

B

[2]

2.6

1.5

2

1.5

6

C

700

340

300

340

300

330

600

1500

400

600

7

N

1

0.5

8

O

330000

346000

246100

420400

323000

420400

318100

230400

360200

367000

101000

388400

9

F

[30]

40

40

20

28

10

Ne

11

Na

6800

7000

7500

4700

7180

5950

7180

5400

3950

5900

1560

1700

6000

12

Mg

138600

139000

135500

95000

159000

123000

159000

120400

87800

134300

136700

38000

135200

13

Al

11000

14000

14300

10000

16100

13800

16100

12200

8800

13900

15200

3900

15870

14

Si

169500

178000

179600

123000

214300

163000

214300

163700

12070

184100

180800

52000

182900

15

P

1300

16000

1500

1800

506

1050

506

800

988

600

1140

1400

520

16

S

20700

20000

20100

10800

20100

21200

20100

33400

18650

19800

18000

4900

19800

17

Cl

[100]

800

470

500

904

1000–1500

904

500

800

18

Ar

19

K

1000

900

900

1200

2630

1540

2630

2000

1430

1700

1740

400

1660

20

Ca

13900

16000

14300

12000

19200

13300

19200

14600

11000

16500

17300

4600

16500

21

Sc

9

5

5

3.5

110

4

5.8

60

22

Ti

800

700

580

600

2100

1320

2100

1150

100

1080

120

23

V

65

80

50

56

300

39

300

165

24

Cr

3000

2500

2700

2200

5000

3430

5000

3900

2900

2800

3220

900

2800

25

Mn

2000

2000

2400

1900

2050

2080

2050

1600

1260

1400

2250

300

1400

26

Fe

286000

256000

241000

457000

127600

288000

127600

299500

471830

243200

233100

720600

213900

27

Co

1000

900

1100

2600

181

1200

181

1300

2250

500

1300

4400

500

28

Ni

16800

14000

14500

35100

2010

15680

20100

19000

35600

13100

16300

65000

13600

29

Cu

100

[40]

170

130

1.55

170

1.55

346

100

140

30

Zn

50

20

76

480

3.4

138

3.4

126

31

Ga

5

8

4.6

20

5

4.2

5

32

Ge

10

[40]

53

82

33

79

33

170

33

As

3

70

18

160

20

20

200

34

Se

9

9

6.7

9

7

46

35

Br

[10]

22

23

1.5

1

20

1

0.55

36

Kr

37

Rb

3

8

8

2.7

4.5

3.5

4.5

2.5

38

Sr

11

22

23

16

72

20

72

40

39

Y

4

5

5.5

4

34

4.72

6.13

8.2

40

Zr

33

90

80

60

100

73

100

55

41

Nb

0.5

0.5

0.41

0.38

2

2

1.5

42

Mo

1.6

5

3.6

8.1

2.5

5.3

2.5

33

43

Tc

44

Ru

4.2

9

45

Rh

0.2

0.6

0.47

1.6

0.8

1.9

46

Pd

[1]

0.5

0.92

1.5

1.53

7.1

47

Ag

0.1

0.5

1.35

1.3

2

2

48

Cd

[0.5]

2

1.6

1.3

4.6

49

In

0.001

0.2

0.2

0.5

0.15

0.044

50

Sn

1

20

14

33

4

20

4

130

51

Sb

0.1

0.4

0.64

0.9

0.1

0.1

1.3

52

Te

1

0.14

0.13

[0.1]

0.93

53

I

0.04

1

1.25

1

1

54

Xe

55

Cs

0.1

0.08

1.1

0.06

0.1

0.08

0.1

0.055

56

Ba

3.4

7

3.9

5.4

20

6.9

20

11

57

La

0.33

0.2

1.9

1.3

1.58

2.05

58

Ce

0.51

2

2.1

1.5

4

1.77?

2.3

2.2

59

Pr

0.12

0.8

0.88

0.6

0.75

0.97

60

Nd

0.63

3

3

2.2

3

2.59

3.37

1.65

61

Pm

62

Sm

0.22

1

1.1

0.78

3

0.95

1.23

1.65

63

Eu

0.083

0.3

0.27

0.2

0.25

0.3

64

Gd

0.34

1.6

1.7

1.2

1.42

1.84

65

Tb

0.37

2

0.52

0.38

0.45

0.58

66

Dy

0.37

2

2.1

1.5

1.8

2.34

67

Ho

0.075

0.6

0.6

0.43

0.51

0.66

68

Er

0.19

1.6

1.7

1.3

1.48

1.93

69

Tm

0.038

0.3

0.31

0.22

0.26

0.34

70

Yb

0.19

1.6

1.7

1.2

1.42

1.85

71

Lu

0.036

0.5

0.54

0.39

0.46

0.6

72

Hf

1.4

0.8

1.6

0.6

1

1.6

1

0.55

73

Ta

0.02

0.3

0.28

0.25

0.7

0.7

0.0077

74

W

0.14

17

16

14

18

15

18

12.6

75

Re

0.08

0.0018

0.08

0.34

0.0008

0.002

0.0008

0.0036

76

Os

0.5

1.1

1.2

3

1.92

3.84

77

Ir

0.5

0.6

0.38

1.2

0.65

0.88

78

Pt

[5]

3

1.9

7.6

0.083

3.25

0.083

6.6

79

Au

0.3

0.26

0.25

0.72

0.65

0.55

80

Hg

[0.1]

[0.009]

<0.01

81

Tl

[0.0004]

[0.14]

0.15

0.016

82

Pb

0.15

[2]

25

5

11

61

83

Bi

[0.003]

[0.16]

0.02

0.2

0.3

90

Th

0.04

0.2

1.2

2

0.8

1.1

92

U

0.014

0.05

2.8

1

0.36

Sum

100.27

101.43

[63.7]

100.68

100.04

97.85

100.00

100.00

100.00

99.88

100.04

99.95

100.00

Note: [Mas62] Mason (1962), mass-ratios assumed were silicate: metal: sulfide = 74.7: 19.6: 5.7, volatile-free. [Lev56] Levin et al. (1956), silicate: metal = 6:1, volatile-free. [U52] Urey (1952), volatile-free chondrite average, for a few elements weighted average of meteoritic silicate: metal: sulfide= 85: 9.5: 5.5, following Goldschmidt, who did not use these observed proportions from Prior (1916). A previous, incorrect version of Urey (1952) appeared in the first edition of Urey’s book The Planets where he used silicate: metal: sulfide= 85: 9.5: 5.5, and Brown’s (1949) “silicates” (= stony meteorites) and “metal (= iron meteorites). Urey had taken the sulfide composition listed by Brown (1949) in addition to the stony and iron meteorites but did not realize that the sulfide portion was already included with the metal portion there. Thus the sulfide component was about twice too high. He corrected his mistake in Urey (1952) and corrected the table in the second edition of his book. [Bro49] Brown (1949), weighted average of stony and iron meteorites to match mantle and core mass of Earth, silicate: metal: sulfide= 100: 67: 0 (67, 33, 0), sulfide automatically determined from sulfides native to stony and iron meteorites. [Budd46] Buddhue (1946), aerolites (excluding iron meteorites). [Gold37] Goldschmidt (1937), weighted average of meteoritic silicate: metal: sulfide = 100: 20: 10 (equivalent to 77: 153: 1.1). [NN34] Noddack and Noddack (1934, 1935), average stony meteorites. [Sas33] Saslawsky (1933), weighted average of meteoritic silicate: metal: sulfide= 100: 25: 7 to match meteorite mean density of 3.34g/cm3. [NN30] Noddack and Noddack (1930), weighted average of meteoritic silicate: metal: sulfide= 100: 68: 9.8 = 82.5: 12: 5.5 to match average density of terrestrial planets taken as 5.10 g/cm3, but see Saslawasky (1933) who noted errors in their mass-balance computation and density assumptions. [Har17] Harkins (1917), average of 125 chondrite analyses. [Mer16] Merrill (1916), average of 53 chondrite analyses. [Far15] Farrington (1915), weighted mixture of stony and iron meteorites 100,257 to match the Earth’s density taken as 5.57 g/cm3. [Mer09] Merrill (1909), average of stony meteorites.

The CI-chondrites are the most volatile element-rich among the known chondritic meteorite groups (Mason, 1960, 1971). The arguments for using CI-chondrites as proxies for the composition of the primordial solar system matter are their good correspondence with elemental abundances, including volatile elements (but not highly volatile elements such as H, C,N,O, noble gases), measured in the solar photosphere, and their continuity in elemental and isotopic abundance trends with mass number. Figure 1 displays the close correspondence of elemental abundances in the Sun and in CI-chondrites.

Figure 1. Atomic abundances of the elements in CI-chondrites versus abundances in the Sun (mainly photosphere). Both abundance sets are normalized to 106 silicon atoms. Perfect agreement would be along a 1:1 line, and most elements plot along such a line within 10% to 20%. Despite some small scatter around the perfect correlation line, the correspondence over 13 orders of magnitude is impressive. Notable exceptions are lithium, which is normal in meteorites but lost from the photosphere by diffusion into the Sun and destroyed there; the other exceptions are elements (the noble gases, H, C, N, and O) that form highly volatile gases not retained in meteorites.

Data from Palme et al. (2014).

Other chondrite groups have elemental fractionations relative to CI-chondrites and are not as good a match for the solar photospheric abundances. Carbonaceous chondrites other than CI-chondrites show clear volatility-related fractionations for some elements (e.g., alkali elements, halogens, In, Tl, Bi) that correlate with condensation temperatures computed for a gas of overall solar composition (see, e.g., Palme et al., 2014). The ordinary and enstatite chondrites also show volatility-related elemental fractionations, and these chondrites further accreted metal and silicate phases in nonsolar proportions as seen from their Si/Fe or Mg/Fe concentration ratios (Si and Mg are typically in silicates, and larger amounts of Fe are typically in the form of metal and sulfide in chondrites). Therefore, these chondrites may not provide good matches to the Sun. However, more recent photospheric abundances measurements using more detailed 3D atmospheric models (see Asplund et al., 2009) gave different results for some diagnostic element ratios such as Si/Fe or Mg/Si, and then the correlation in Figure 1 would show more scatter. However, the basic feature remains—a correlation over 13 orders of magnitude for all elements except the most volatile ones and lithium. In addition, there are uncertainties in the photospheric abundances (typically 7% [1-sigma] or higher) and the CI-chondrite abundances (typically 3%–5% [1-sigma] or higher), which must be considered for the comparison. As described later, measurements of the solar wind from the Genesis mission can provide high-level accuracy measurements for some major elements, and initial results seem to indicate better agreement with CI-chondrites than with the spectroscopic 3D photospheric results (e.g., Burnett et al., 2017). The use of CI-chondrites as abundance standard started during the 1950s but was questioned in the interim, and ample literature about this issue exists. The earlier analysts of CI-chondrites—Berzelius in 1835 investigating the Alais CI-chondrite (fall 1806) and Daubree (1864) the CI-chondrite Orgueil (fall 1864)—already suspected these rare carbonaceous chondrites could be close to “primeval matter” mainly because of their relative high contents of organic compounds. Urey (1952) concluded the carbonaceous chondrites were the best representatives of nonvolatile solar system material because they contain abundant water, S, C, alkalies, and other volatiles, but he (Urey & Edwards, 1956) abandoned this idea because the carbonaceous chondrites showed more variation in volatile element contents from his modern analyses of alkali elements than he expected from assumed uniform primordial condensations. Then Urey favored ordinary chondrites for a brief time. However, as Ringwood (1961) remarks, Urey used compositional averages of all different types of carbonaceous chondrites instead of only CI-chondrites to conclude this. Ordinary chondrites can have higher alkali element contents than some types of carbonaceous chondrites, but this is not true for larger samples of CI-chondrites.

The use of CI-chondrites became firmly established after measurements by Reed, Kigoshi, and Turekevitch (1960) for the volatile elements Bi, Pb, and Tl in CI-chondrites provided a much better match to the theoretical solar system abundances estimated by Suess and Urey (1956) and Cameron (1959) from nuclear systematics. In his subsequent abundance tables Cameron (1968, 1973) made extensive use of CI-chondrite data and set the use of CI-chondrites into motion.

It took until the 1970s to completely settle this issue because there was a severe disagreement between the CI-chondritic (and other chondrites) and the photospheric iron abundance (e.g., Anders, 1971; Withbroe, 1971). Abundant elements such as Mg and Si agreed reasonably well, but the solar iron value appeared to be a factor of 10 times lower than in meteorites. Further, the photospheric and coronal Fe abundances were discordant, and the discovery of the forbidden transition lines of singly ionized Fe (written as [Fe II]) in the photosphere hinted at an iron abundance an order of magnitude larger than previously thought (e.g., Grevesse & Swings, 1969; Nussbaumer & Swings, 1970). Eventually the atomic parameters (the gf-values) for iron measured in the laboratory were found to be off by a factor of 10 because the temperature calibration of the furnace used in the measurements was incorrect. New laboratory oscillator strengths for permitted neutral (Fe I) and ionized (Fe II) iron lines were measured, and the Fe abundance was redetermined (Baschek, Garz, Holweger, & Richter, 1970; Garz, Holweger, Kock, & Richter, 1969; Garz & Kock, 1969; Garz, Kock, et al., 1969). The abundances of Fe from the Fe I and Fe II lines then agreed far better with values from the forbidden [Fe II] lines, the corona, and meteorites. This solidified the use of CI-chondritic compositions as the solar system standard of condensable elemental abundances.

There are five known CI-chondrite falls (Alais, Ivuna, Orgueil, Tonk, Revelstoke), and all are of petrologic type 1. (For a while the Tagish Lake meteorite fall was thought to be a CI2 chondrite, but this was incorrect). As noted earlier, the petrological type 1 designates completely aqueously altered meteorites, and the mineralogy of CI-chondrites is not that of pristine condensates from the solar nebula. Their mineralogy was changed on their parent asteroid after condensates accreted, but, importantly, their overall elemental composition was not changed. Thus the mineral alteration occurred in a closed system without element additions or losses (isochemical alteration).

Abundances of the Elements in CI-Chondrites

The elemental composition by mass of CI-chondrites from various sources with somewhat different approaches is given in Table 3. In 1864, the Orgueil CI-chondrite fell in France, and about 12 kg are left of it in meteorite collections. This is more than 90% of all mass preserved from all five known CI-chondrite falls. The number of analyses of Orgueil dominates the abundance statistics, and several authors mainly use only Orgueil data (indicated by “Orgueil” in Table 3). Others use straight or weighted compositions of all five CI-chondrites to calculate the group meteorite composition (e.g., Lodders, 2003; Wasson & Kallemeyn, 1988). For some elements, some authors prefer to derive concentrations of one element from element concentration ratios measured in other meteorite groups. This approach exploits the chemical similarities (leading to often constant concentration ratios) of elements during fractionation processes in meteorites. Then abundances can be computed from well-determined ratios in other chondrites if an element in the ratio is hard to measure in CI-chondrites (see Anders & Ebihara, 1982; Anders & Grevesse, 1989; Lodders et al., 2009; Palme et al., 2014). The “ratio approach” was often simply a necessity because good analytical data for CI-chondrites were sparse (only five meteorites, of which only two to three are available in sufficient quantity), and, for several elements, the number of reliable analyses remains precariously sparse (see Lodders, 2003).

An often unappreciated complication for analyses of CI-chondrites is their high water content (around 15%–20%) in the form of crystal water (in epsomite) and in hydrated silicates. Sometimes analytical results are reported for dried samples, so reported absolute element concentrations will be larger than in the native meteorite. In other instances, analytical results are reported for devolatilized samples, which were water-free and carbon- and nitrogen-free (carbonates, N-bearing organics, free carbon, and ammonia were driven off by heating, and, in some instances, all or some portions of sulfur were also removed). Therefore, individual analyses must be checked for sample processing before they can be included into evaluations of mean CI-chondrite compositions. Older average meteorite compositions listed in Table 2 sometimes refer to “volatile-free” compositions, which is indicated in the table footnote.

Table 3. Elemental Compositions by Mass of CI-Chondrites

Z

This Review

2σ

Orgueil [PLJ14]

Orgueil [LPG09]

Orgueil [PJ03]

Orgueil [L03]

all CI [L03]

Orgueil [MS95]

Orgueil [PB93]

Orgueil [AG89]

all CI [WK88]

Orgueil [P88]

all CI [W85]

All CI [AE82]

Orgueil [PSZ81]

all CI [Yav68]

1

H

18600

3440

19700

19700

20200

19700

21015

20200

20200

20000

20000

20200

2

He

0.00917

0.00917

0.00917

0.00917

0.00917

0.00917

0.00917

0.018

0.00917

3

Li

1.51

0.12

1.45

1.47

1.49

1.45

1.46

1.5

1.49

1.5

1.57

1.45

1.57

1.59

1.45

1.3

4

Be

0.0220

0.0016

0.0219

0.0210

0.0249

0.019

0.0252

0.025

0.0249

0.0249

0.027

0.025

0.027

0.027

0.025

5

B

0.744

0.172

0.775

0.775

0.69

0.775

0.710

0.9

0.87

0.87

1.2

0.27

1

1.25

0.27

6

C

41300

8400

34800

34800

32200

34800

35180

35000

32200

34500

32000

35000

32000

34500

35000

7

N

2500

660

2950

2950

3180

2950

2940

3180

3180

3180

1500

3180

1500

3180

8

O

453840

20000

459000

458500

465000

462300

458190

465000

464000

460000

470000

460000

464000

470000

9

F

92

40

58.2

58.2

58.2

58.2

60.6

60

58.2

60.7

64

54

64

58.2

54

10

Ne

1.80E-4

1.80E-4

1.80E-4

1.68E-04

1.80E-4

1.80E-4

1.68E-4

3.0E-04

1.68E-4

11

Na

5100

500

4962

4990

4982

4990

5010

5100

4982

5000

4900

5020

4800

4830

5020

5500

12

Mg

95170

4000

95400

95800

96100

96000

95870

96500

96100

98900

97000

94000

97000

95500

93600

96000

13

Al

8370

600

8400

8500

8490

8580

8500

8600

8650

8680

8600

8200

8600

8620

8200

8600

14

Si

107740

7200

107000

107000

106800

105700

106500

106500

106800

106400

105000

107000

105000

106700

106800

106000

15

P

978

120

985

967

926

924

920

1080

1105

1220

1020

1010

1020

1180

1010

16

S

53600

4400

53500

54100

54100

53500

54100

54000

52500

62500

59000

58000

59000

52500

58000

61300

17

Cl

717

220

698

698

698

700

704

680

698

704

680

678

680

698

678

18

Ar

0.00133

0.00133

0.00133

0.00123

0.00133

0.00133

0.00123

0.0013

0.00123

19

K

539

48

546

544

544

543

531

550

544

558

560

517

560

569

517

560

20

Ca

8840

700

9110

9220

9320

9080

9070

9250

9510

9280

9200

9000

9200

9020

9000

11100

21

Sc

5.83

0.40

5.81

5.90

5.9

5.84

5.83

5.92

5.9

5.82

5.8

5.9

5.8

5.76

5.9

5.6

22

Ti

450

30

447

451

458

439

440

440

441

436

420

440

420

436

440

420

23

V

53.6

4.0

54.6

54.3

54.3

56.0

55.7

56

54.3

56.5

55

55.6

56

56.7

55.6

24

Cr

2610

200

2623

2650

2646

2630

2590

2650

2646

2660

2650

2670

2650

2650

2670

2400

25

Mn

1896

160

1916

1930

1933

1920

1910

1920

1933

1990

1900

1820

1900

1960

1820

1600

26

Fe

185620

13000

186600

185000

184300

183500

182800

181000

182300

190400

182000

183000

182000

185100

183000

27

Co

508

30

513

506

506

507

502

500

506

502

508

501

508

509

501

28

Ni

10950

700

10910

10800

10770

10670

10640

10500

10770

11000

10700

10800

10700

11000

10800

29

Cu

130

20

133

131

131

131

127

120

131

126

121

108

121

112

108

30

Zn

311

20

309

323

323

318

310

310

323

312

312

347

312

308

347

31

Ga

9.45

0.70

9.62

9.71

9.71

9.57

9.51

9.2

9.71

10

9.8

9.1

9.8

10.1

9.1

32

Ge

33.4

3.0

32.6

32.6

32.6

33.1

33.2

31

32.6

32.7

33

31.3

33

32.2

31.3

33

As

1.77

0.16

1.74

1.74

1.81

1.70

1.73

1.85

1.81

1.86

1.84

1.85

1.84

1.91

1.85

34

Se

20.4

1.6

20.3

20.7

21.4

19.5

19.7

21

21.3

18.6

19.6

18.9

19.6

18.2

19.9

35

Br

3.77

1.80

3.26

3.26

3.5

3.18

3.43

3.57

3.5

3.57

3.6

3.56

3.6

3.56

2.53

36

Kr

5.22E-5

5.22E-5

5.22E-5

2.98E-5

5.22E-5

5.22E-5

2.98E-5

3.3E-5

2.98E-5

37

Rb

2.22

0.18

2.32

2.31

2.32

2.12

2.13

2.3

2.32

2.3

2.22

2.06

2.22

2.30

2.06

2.3

38

Sr

7.79

0.50

7.79

7.81

7.26

7.67

7.74

7.25

7.26

7.8

7.9

8.6

7.9

7.91

8.6

<1

39

Y

1.50

0.10

1.46

1.53

1.56

1.6

1.53

1.57

1.57

1.56

1.44

1.57

1.44

1.50

1.44

1.6

40

Zr

3.79

0.28

3.63

3.62

3.86

4.02

3.96

3.82

3.87

3.94

3.8

3.87

3.8

3.69

3.82

11

41

Nb

0.279

0.015

0.283

0.279

0.247

0.259

0.265

0.24

0.246

0.246

0.27

0.246

0.27

0.250

0.3

42

Mo

0.976

0.050

0.961

0.973

0.928

0.929

1.02

0.9

0.928

0.928

0.92

0.92

0.92

0.92

0.92

43

Tc

44

Ru

0.666

0.04

0.690

0.686

0.683

0.725

0.692

0.71

0.714

0.712

0.71

0.69

0.71

0.714

0.69

45

Rh

0.133

0.008

0.132

0.139

0.14

0.139

0.141

0.13

0.134

0.134

0.134

0.13

0.134

0.134

0.13

46

Pd

0.558

0.030

0.560

0.558

0.556

0.600

0.588

0.55

0.556

0.56

0.56

0.53

0.56

0.557

0.53

47

Ag

0.204

0.008

0.201

0.201

0.197

0.203

0.201

0.2

0.197

0.199

0.208

0.21

0.208

0.220

0.21

48

Cd

0.679

0.024

0.674

0.674

0.68

0.675

0.675

0.71

0.68

0.686

0.65

0.77

0.65

0.673

0.77

49

In

0.0786

0.0040

0.0778

0.0778

0.078

0.0779

0.0788

0.08

0.0778

0.08

0.08

0.08

0.08

0.0778

0.08

50

Sn

1.63

0.16

1.63

1.63

1.68

1.66

1.68

1.65

1.68

1.72

1.72

1.75

1.72

1.68

1.75

51

Sb

0.169

0.018

0.145

0.149

0.133

0.148

0.152

0.14

0.133

0.142

0.153

0.13

0.153

0.155

0.13

52

Te

2.31

0.18

2.28

2.28

2.27

2.26

2.33

2.33

2.27

2.32

2.4

2.34

2.4

2.28

2.34

53

I

0.77

0.62

0.53

0.53

0.43

0.48

0.480

0.45

0.433

0.433

0.5

0.56

0.5

0.43

0.56

54

Xe

1.74E-4

1.74E-4

1.74E-4

4.62E-05

1.74E-4

1.74E-4

4.62E-5

3.2E-5

4.62E-5

55

Cs

0.188

0.012

0.188

0.188

0.188

0.184

0.185

0.19

0.188

0.187

0.183

0.19

0.183

0.186

0.19

0.18

56

Ba

2.39

0.16

2.42

2.41

2.41

2.30

2.31

2.41

2.41

2.34

2.3

2.6

2.3

2.27

2.2

2.4

57

La

0.244

0.016

0.241

0.242

0.245

0.238

0.232

0.237

0.245

0.2347

0.236

0.245

0.236

0.236

0.245

0.19

58

Ce

0.627

0.052

0.619

0.622

0.638

0.637

0.621

0.613

0.638

0.6032

0.616

0.638

0.616

0.619

0.638

0.62

59

Pr

0.0951

0.0066

0.0939

0.0946

0.0964

0.0925

0.0928

0.0928

0.0964

0.0891

0.0929

0.096

0.0929

0.090

0.096

0.09

60

Nd

0.472

0.036

0.474

0.471

0.474

0.459

0.457

0.457

0.474

0.4524

0.457

0.474

0.457

0.462

0.474

0.42

61

Pm

...

62

Sm

0.153

0.012

0.154

0.152

0.154

0.145

0.145

0.148

0.154

0.1471

0.149

0.154

0.149

0.142

0.154

0.13

63

Eu

0.0577

0.0050

0.0588

0.0578

0.058

0.0548

0.0546

0.0563

0.058

0.056

0.056

0.058

0.056

0.0543

0.058

0.053

64

Gd

0.208

0.018

0.207

0.205

0.204

0.201

0.198

0.199

0.204

0.1966

0.197

0.204

0.197

0.196

0.204

0.24

65

Tb

0.0380

0.0030

0.0380

0.0384

0.0375

0.0372

0.0356

0.0361

0.0375

0.0363

0.0355

0.037

0.0355

0.0353

0.037

0.044

66

Dy

0.252

0.020

0.256

0.255

0.254

0.239

0.238

0.246

0.254

0.2427

0.245

0.254

0.245

0.242

0.254

0.22

67

Ho

0.0563

0.0044

0.0564

0.0572

0.0567

0.0555

0.0562

0.0546

0.0567

0.0556

0.0547

0.057

0.0547

0.054

0.057

0.056

68

Er

0.164

0.012

0.166

0.163

0.166

0.167

0.162

0.16

0.166

0.1589

0.16

0.166

0.159

0.160

0.166

0.14

69

Tm

0.0259

0.0024

0.0261

0.0261

0.0256

0.0245

0.0237

0.0247

0.0256

0.0242

0.0247

0.026

0.0247

0.022

0.026

0.022

70

Yb

0.167

0.014

0.169

0.169

0.165

0.165

0.163

0.161

0.165

0.1625

0.159

0.165

0.159

0.166

0.165

0.13

71

Lu

0.0249

0.0020

0.0250

0.0253

0.0254

0.0235

0.0237

0.0246

0.0254

0.0243

0.0245

0.025

0.0245

0.0243

0.025

0.023

72

Hf

0.106

0.008

0.107

0.106

0.107

0.116

0.115

0.103

0.107

0.104

0.12

0.106

0.12

0.119

0.12

73

Ta

0.0148

0.0014

0.015

0.0145

0.0142

0.0145

0.0144

0.0136

0.014

0.0142

0.016

0.014

0.016

0.017

0.014

74

W

0.102

0.014

0.096

0.0960

0.0903

0.089

0.089

0.093

0.095

0.0926

0.1

0.093

0.1

0.089

0.089

75

Re

0.0369

0.0028

0.040

0.0393

0.0395

0.0379

0.0371

0.04

0.0383

0.0365

0.037

0.037

0.037

0.037

0.037

76

Os

0.475

0.020

0.495

0.493

0.506

0.487

0.486

0.49

0.486

0.486

0.49

0.49

0.49

0.699

0.49

77

Ir

0.474

0.020

0.469

0.469

0.48

0.470

0.470

0.455

0.459

0.481

0.46

0.48

0.46

0.473

0.48

78

Pt

0.931

0.072

0.925

0.947

0.982

0.962

1.00

1.01

0.994

0.999

0.99

1.05

0.99

0.953

1.05

79

Au

0.147

0.024

0.148

0.146

0.148

0.145

0.146

0.14

0.152

0.14

0.144

0.14

0.144

0.145

0.14

80

Hg

0.288

0.140

0.35

0.350

0.31

0.345

0.314

0.3

0.31

0.258

0.39

(5.3)

0.39

0.39

(5.3)

81

Tl

0.141

0.014

0.140

0.142

0.143

0.144

0.143

0.14

0.143

0.142

0.142

0.14

0.142

0.143

0.14

82

Pb

2.64

0.16

2.62

2.63

2.53

2.58

2.56

2.47

2.53

2.47

2.4

2.43

2.4

2.43

2.43

83

Bi

0.113

0.016

0.110

0.110

0.111

0.108

0.110

0.11

0.111

0.114

0.11

0.11

0.11

0.111

0.11

90

Th

0.0298

0.0030

0.0300

0.0310

0.0298

0.0312

0.0309

0.029

0.0298

0.0294

0.029

0.029

0.029

0.0286

0.029

0.065

92

U

0.00816

0.00106

0.0081

0.0081

0.0078

0.0081

0.0084

0.0074

0.0078

0.0081

0.0082

0.0082

0.0082

0.0081

0.0082

0.024

Sum

%

100.00

100.07

99.98

100.38

100.00

99.81

[51.81]

100.07

102.28

99.83

98.77

99.82

100.38

98.77

[29.4]

Note: This review: Weighted average of all CI-chondrites as done in Lodders (2003). Uncertainties are either the deviation of the weighted group mean or the standard deviation found for individual meteorites, whichever is larger. The quoted uncertainties of CI-chondritic abundances thus should reflect upper limits to the characteristic variations in elemental abundances obtained on different samples and by different analytical methods.

Sources: [AE82] Anders and Ebihara (1982); [AG89] Anders and Grevesse (1989); [L03] Lodders (2003); [LPG09] Lodders et al. (2009); [MS95] McDonough and Sun (1995); [PJ03] Palme and Jones (2003); [PLJ14] Palme et al. (2014); [PSZ81] Palme, Suess, and Zeh (1981); [W85] Wasson (1985); [WK88] Wasson and Kallemeyn (1988); [Yav68] Yavnel (1968).

The CI-chondrite compositions by mass are easily converted into atomic abundances, which traditionally is done by setting the number of silicon to 106 atoms. All data are converted using the atomic weights calculated for the adopted isotopic composition of each element using the individual isotope masses from Pfeiffer, Venkataramaniahc, Czoka, and Scheidenberger (2014). Elements can have different solar isotopic compositions than on the Earth, so the atomic weights will be different than those listed for terrestrial materials (e.g., C, N, O, noble gases). For example, the atomic weight of Ar in air on Earth is 39.948 amu (atomic mass units where 12C=12.0) and is dominated by the mass of abundant 40Ar. For solar, the relative masses of 36Ar and 38Ar dominate to give 36.275 amu because 40Ar is extremely rare in the solar wind.

Atomic abundance ratios have been used since the 1930s, and the use of atom percentages is not deemed practical for modeling solar system abundances. Using a relative atomic scale circumvents the problems with compositions reported on a “volatile-free” basis since the atomic ratios of refractory elements relative to silicon are unaffected. Data from several studies listed in Tables 2 and 3 converted to the atomic scale relative to silicon are used in Table 6.

Solar Abundances of the Elements, Photosphere, Sunspots, Corona, and Solar Wind

As mentioned earlier, the Sun contains more than 99 mass percent of solar system materials, and therefore the Sun’s composition should represent a good average composition of the elements within the entire solar system. There is one notable exception: After the Sun formed from the proto-solar nebula and before hydrogen fusion began, the young Sun was fully convective and was burning deuterium during its T-Tauri stage, which is the reason for the absence of deuterium in the solar photosphere. The current Sun is not uniform in composition and has a radial abundance gradient. The Sun is a main sequence star, and it operates by nuclear fusion of hydrogen in its center. The hydrogen-burning reactions through the so-called pp-fusion chain (ca 97%) and CNO-burning cycle (ca. 3%) steadily decrease the amount of hydrogen and increase the amounts of helium in the center. The interior of the Sun is not fully convective, and energy is transported from the core outward through radiation. Only the upper envelope of the Sun is now convective (see the review by Hanasoge, Gizon, & Sreenivasan, 2016) and is believed to be relatively homogeneous in composition. The convective envelope (a shell about one-third of the Sun’s total radius), and because the convective layer only has minimal exchange of matter with the central Sun, the photosphere near the top of the convective zone should give a good representative composition for the Sun (and therefore the solar system) at the time of its formation.

Over the Sun’s lifetime, some helium and all other heavy elements gravitationally settled from the convective zone into the solar interior. Thus the photosphere (near the top of the convection zone) no longer represents the original proto-solar composition. Gravitational settling is in part counteracted by radiation levitation, which acts differently on different ions (Eddington, 1926). These effects were investigated by, for example, Turcotte, Richter, Michaud, Iglesia, and Rogers (1998), Vauclair (1998), Turcotte and Wimmer-Schweingruber, (2002), and Yang (2016, 2019). Yang (2019) took the effects of rotation on mixing into account (see later discussion). In the present-day convective envelope, the amounts of He and the heavier elements are about 10% to 20% smaller than at the time the Sun formed (depending on models; see later discussion). However, except for Li, the relative abundances of the elements heavier than helium are assumed to be more or less unchanged, and any potential differences to meteoritic values as a function of mass are not resolvable with current abundance data. Therefore heavy element abundance ratios are taken as approximately constant (e.g., Asplund et al., 2009; Lodders et al., 2009; Piersanti et al., 2007). The notable exception is lithium. The lithium relative abundance in the solar photosphere is 170 times lower than that in meteorites. This is due to pre-mainsequence lithium destruction and ongoing settling plus the nuclear destruction of the fragile lithium nuclei at the hot bottom of the solar convection zone.

The most reliable abundances for the present-day Sun are obtained from photospheric absorption spectra. Even today not all elements can be detected or determined quantitatively in the Sun because the solar spectrum is crowded with thousands of lines (e.g., neutral Fe and ionized Fe lines are prolific). Crowded spectra complicate the spectral analyses because of line blending, which makes the detection of lines from elements with low excitation potentials difficult because their absorption lines are very weak. Some elements with low abundances may only have few useful spectral lines at wavelengths difficult to access (e.g., in the ultraviolet [UV] part of spectrum that is not accessible through the Earth’s atmosphere). Other elements (e.g., F, Cl, Tl) can only be measured from the spectra of the cooler sunspots where many more atomic and molecular lines complicate the spectra. For some elements (As, Se, Br, Te, I, Cs, Ta, Re, Hg, Bi, Th, noble gases), no or only unreliable photospheric determinations exist because there are no observable or only heavily blended lines in the solar spectrum. The noble gases He, Ne, Ar, Kr, and Xe have no lines in the photospheric spectrum, and their abundances could only be derived from observations of coronal sources such the solar wind, solar flares, or solar energetic particles (SEPs).

The He abundance is derived indirectly by solar models in conjunction with observations from helioseismology. The solar surface shows up- and downward oscillations caused by some sound waves created in the turbulent convection zone and reflected at the photospheric region (“the solar surface”). Helioseismology can probe the entire structure of the Sun through observations of the different sound modes, and because sound speed is dependent on the composition of the medium that sound passes through, constraints on composition (especially the mass fractions of X, Y, and Z noted earlier) are obtained. However, matching the seismic observations (and to derive the helium abundance, which is a free parameter in the solar models) depends on the details of the standard solar models. The standard models must explain how the Sun works over time (solar radius, density structure, neutrino flux, depth of convection zone; see, e.g., Basu & Antia, 2004, 2008; Basu, Grevesse, Mathis, & Turck-Chieze, 2015; Serenelli & Basu, 2010, Serenelli, Haxton, & Pena-Gray, 2011). The overall solar composition derived from the elements in the photosphere is extremely important for the standard solar model because the abundant elements C, N, O, Ne, Fe, and lesser abundant heavy elements determine the opacity of the solar interior, which in turn determines how energy is transported outward (radiation, convection). The C, N, and O abundances have been revised substantially downward since the early 2000s. The elements C, N, and O are about two-thirds of all heavy elements by mass. The lower abundances gave lower opacities that conflicted with those required from helioseismology. A surge of new measurements and theoretical work has not yet resolved all disagreements (see, e.g., Basu et al., 2015; Serenelli & Basu, 2010; Vinyoles et al., 2017; Yang, 2016, 2019).

The helium abundance in Table 4 for the recommended 3D and 1D photospheric compositions is obtained as follows. The photospheric abundances relative to hydrogen for all other elements can be independently converted from atomic abundances to mass abundances, which gives the mass fraction ratio for all elements heavier than helium relative to the hydrogen mass as Z/X=0.0149 for the recommended 3D photospheric values in Table 4. Basu and Antia (2004) derived the hydrogen mass fraction X in the solar convection zone from the observed solar oscillation frequencies by the Global Oscillations Network Group (Hill et al., 1996) and by the Michelson Doppler Imager (Schou et al., 1998). They found a hydrogen mass fraction X=0.7389, which was insensitive to the mass fraction ratio Z/X assumed in their two calibration models (with Z/X=0.0171 and 0.0218). The Y/X values for the recommended abundances here are between the values used in the calibration models by Basu and Antia. Using their determined mass fraction X, the mass fraction of helium is then easily obtained as Y=0.2462 (since Y/X=1/X1Z/X and X+Y+Z=1). Once the Y/X ratio is converted back to an atomic ratio and properly scaled, the photospheric helium abundance is A(He)=10.924 (the 1D abundances with X/Z=0.0153 give X/Z=0.0153 and Y=0.2458). For more related work on the solar mass fractions, see the section on solar system abundances.

Photospheric Abundances

The quantitative determination of photospheric abundances involves construction of a numerical model atmosphere, calculation of the emitted spectrum based on the model atmosphere, and comparison of this spectrum with the observed spectrum. A comprehensive review about techniques, problems, and results from solar and stellar spectroscopy is given by Allende-Prieto (2016), and the reader is encouraged to consult this review for reference.

The photosphere is the top layer above the solar convection zone in which the solar gas is well mixed. Suitable atmospheric models describing the temperature, densities, total pressure, composition, and convective mixing as a function of depth were developed to describe the spectral line formation regions in the photosphere and sunspots. Strömgren (1940) constructed the first photospheric model atmosphere with temperature, density, and electron pressure as function of optical depth after Wildt (1939) identified negatively charged hydrogen (H) as the cause of the continuous opacity in the solar atmosphere. This model allowed for the first time the determination of elements relative to hydrogen because the equivalent widths of the observed absorption lines in the spectra depend on the continuum absorption from hydride ions. Unsöld (1948) and Claas (1951) did the first detailed abundance determinations using model atmospheres. Model atmospheres evolved from 1D hydrostatic models (e.g., Holweger, 2001; Holweger & Müller, 1974) to 3D models and can take time-dependent hydrodynamics and granulation into account (Asplund et al., 2009; Caffau, Ludwig, Steffen, Freytag, & Bonifacio, 2011). It is still debated if 3D models produce more accurate solar abundances than 1D models. Differences among various 3D models are apparently small if the same line selections and local thermodynamic equilibrium (LTE) corrections are applied (see later discussion), so line selections, number of lines included in the analyses, and atomic properties remain major issues when 3D results are compared.

Emission from the deeper regions of the photosphere is absorbed in overlying layers and leads to the absorption lines in photospheric spectra. The models of spectral line formation require atomic physics parameters such as accurate wavelengths, ionization potentials, atomic term structures, partition functions, oscillator strengths, transition probabilities and lifetimes, and hyperfine splitting, which must be determined experimentally, especially for the heavy elements because quantum mechanical computations become extremely complex. Goldberg, Müller, and Aller (1960) first applied experimental oscillator strengths to the absorption lines and could obtain refined abundances for 42 elements. Much progress has been made, and laboratory measurements of transition probabilities and lifetimes are ongoing and remain an active research area in astronomy (e.g., see, Lawler et al., 2019, 2017, 2014, 2015; Snenden et al., 2016, 2009; and references therein), and solar abundance determinations will continue to improve.

An important consideration in solar abundance determinations is departures from LTE, that is, whether “the quantum-mechanical states of atoms, ions, and molecules are populated according to the convenient relations of Boltzmann and Saha, valid strictly in thermodynamic equilibrium” (see Holweger, 2001, p. 23; Allende-Prieto, 2016). The assumption of LTE may not be justified for a highly inhomogeneous and dynamic plasma permeated by an intense, anisotropic radiation. Consequences of deviations from LTE are changes in line intensities and in the line shapes of strong lines, and these effects are apparent when observations from the center of the solar disc are compared to those from the limb. Newer spectrum synthesis calculations take effects of non-local thermodynamic equilibrium (NLTE) more regularly into account, which, however, are computationally intensive.

Table 4 summarizes photospheric abundance compilations from Russell (1929) onward. The original tables should be consulted for the measurement details and/or sources of compiled data. The photospheric abundances have changed considerably since the abundance reviews by Anders and Grevesse (1989) and Grevesse and Sauval (1998), largely due to the application of 3D model atmospheres and NLTE corrections. Downward revisions of elemental abundances started with the photospheric oxygen abundance (Allende Pietro et al., 2001), and other elements followed as 3D model atmospheres were applied in solar spectrum analyses (see, e.g., Asplund, Grevesse, & Sauval, 2005; Asplund et al., 2009; Caffau et al., 2011). Differences between different 3D models are ascribed to different selections of lines and atomic parameters as evaluated by different groups. Another change is the more frequent application of NLTE corrections to 1D and 3D analyses to account for collisional line excitations. The problems arising from the lower new solar “3D abundances” of oxygen and other abundant highly volatile elements have become known as the “oxygen crisis,” and the well-working standard solar models have an “opacity problem” (e.g., Basu & Antia, 2005, 2008; Vinyoles et al., 2017; and later discussion here).

Asplund et al. (2009) presented self-consistent photospheric abundances based on their 3D model atmospheres with much improved, realistic physics and what they regarded as the best-available atomic parameters and line selections. Scott, Asplund, et al. (2015), Scott, Grevesse, et al. (2015), and Grevesse, Scott, Asplund, and Sauval (2015) subsequently gave the complete descriptions of the analyses by Asplund et al.; these are also shown in Table 4.

Palme et al. (2014) did not adopt the photospheric abundances based on 3D models by Asplund et al. (2009) because at the time the detailed descriptions for the Asplund et al. solar abundances were pending. Palme et al. compiled photospheric abundances from different papers (as in Lodders, 2003; Lodders et al., 2009). This approach relies on the recommendations given in various papers, and sometimes averages of several studies were used. This may not have a self-consistent base for atmospheric models, since 3D results from different groups, as well as 1D and 3D model results, were included. The recommended photospheric abundances here are listed for “1D” and “3D” separately in Table 4, which also illustrates the generally lower photospheric abundances from 3D models.

The photospheric 3D abundances from Asplund et al. (2009) of several rock-forming elements compare somewhat worse to CI-chondrites (Lodders, 2019) and to results from the solar wind Genesis measurements whereas the latter correspond better to CI-chondrites (e.g., Burnett et al., 2017; Heber, McKeegan, Bochsler, Duprat, & Burnett, 2014). Figure 2 illustrates this situation for Mg, Si, and Fe, which are the three most abundant rock-forming elements (after oxygen).

Figure 2. Major element ratios in CI-chondrites (top row) and from analyses of the solar photosphere. The solar wind Mg/Fe datum is from Genesis data; see Table 5 for references. The block of photospheric abundances is in reverse chronological order. Near the top is for 3D LTE (Asplund et al., 2005), and for 3D NLTE results (Asplund et al., 2009), revised/updated values (Scott, Asplund, et al., 2015; Scott, Grevesse, et al., 2015), and further updates from Armasi and Asplund (2017), Bergemann et al. (2017), and Lind et al. (2017). Scott, Grevesse, et al. also reported 1D results for comparison. The lower block shows the ratios from earlier compilations which includes many analyses with 1D model atmosphere and limited NLTE corrections (see footnote to Table 4 for abbreviations of references used). The error bars of the ratios are calculated by quadratic combination of the elemental uncertainties reported in the references. The grey bar shows the 1-sigma uncertainties of the CI-chondrite ratios from the recommended values in Table 3.

The top row in Figure 2 shows Mg/Si, Fe/Si and Mg/Fe ratios in CI-chondrites and the Mg/Fe ratio for the solar wind. The lower block of data in Figure 2 are photospheric abundances mainly from 1D and LTE analyses and show that the Mg/Si, Fe/Si, and Mg/Fe ratios seem to agree better with respective ratios in CI-chondrites; this includes the 3D LTE results from Asplund et al. (2005). (The Fe/Si and Fe/Mg from Anders & Grevesse [1989] used an incorrect Fe value; see previous discussion on the Fe abundance controversy in meteorites and the Sun.) However, 3D NLTE abundances by Asplund et al. (2009) deviate more from the CI-chondritic ratios, and their Mg/Si ratio is higher than in CI-chondrites or any chondrite group. Subsequent updates to the solar 3D NLTE abundances seem to shift the abundance ratios closer to the CI-chondritic values. The agreement of major element abundances between the solar photosphere and CI-chondrites was one justification to use CI-chondrites as an abundance standard and to use CI-chondrite abundances for elements that cannot be determined for the Sun. If real compositional differences exist between the photosphere, the solar wind, and CI-chondrites that are not already understood within the framework of current models, approaches to derive a representative solar system composition will have to be revised. More work on photospheric abundances is needed to understand differences for other elements between meteoritic and 3D NLTE photospheric abundances.

Table 4. Elemental Abundances in the Solar Photospherea Part 1.

Z

Recommended 3D

Recommended 1D

[Z16] 1D NLTE

[SL+] 1D

[SSG15] 3D NLTE

[SSG15] 1D NLTE

[PLJ14]

[C11] 3D NLTE

[LPG09]

[AGSS09] 3D

[GAS07]

[AGS05]

[L03]

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

1

H

12

12

12

12

12

12

12

12

12

12

12

12

12

2

He

[10.924]in

0.02

[10.923]in

0.02

10.925

0.02

10.925

0.02

10.930

0.01

10.93

0.01

10.93

0.01

10.899

0.01

3

Li

1.04

0.10

1.04

0.10

1.03

0.03

1.03

0.03

1.10

0.10

1.05

0.10

1.05

0.10

1.05

0.10

1.10

0.10

4

Be

1.38

0.09

1.38

0.09

1.38

0.09

1.38

0.09

1.38

0.09

1.38

0.09

1.38

0.09

1.15

0.20

5

B

2.70

0.20

2.70

0.20

2.70

0.17

2.70

0.17

2.70

0.20

2.70

0.20

2.70

0.20

2.70

0.21

6

C

8.47

0.06

8.58

0.11

8.43

0.02

8.50

0.06

8.50

0.06

8.39

0.04

8.43

0.05

8.39

0.05

8.39

0.05

8.39

0.04

7

N

7.85

0.12

8.00

0.11

7.86

0.12

7.86

0.12

7.86

0.12

7.83

0.05

7.78

0.06

7.78

0.06

7.83

0.11

8

O

8.71

0.04

8.76

0.06

8.78

0.04

8.73

0.07

8.76

0.07

8.73

0.07

8.69

0.03

8.66

0.05

8.66

0.05

8.69

0.05

9

F

[4.40]sp

0.25

[4.40]sp

0.25

4.56

0.30

4.56

0.30

4.56

0.30

4.56

0.30

4.56

0.30

4.56

0.30

10

Ne

[8.15]in

0.10

[8.10]in

0.10

8.05

0.10

8.05

0.10

7.93

0.10

7.84

0.06

7.84

0.06

7.87

0.10

11

Na

6.21

0.04

6.28

0.04

6.29

0.02

6.21

0.04

6.27

0.04

6.30

0.03

6.30

0.03

6.24

0.04

6.17

0.04

6.17

0.04

6.30

0.03

12

Mg

7.56

0.05

7.50

0.05

7.51

0.05

7.59

0.04

7.63

0.04

7.54

0.06

7.54

0.06

7.60

0.03

7.53

0.09

7.53

0.09

7.54

0.06

13

Al

6.43

0.04

6.41

0.03

6.39

0.02

6.43

0.04

6.48

0.04

6.47

0.07

6.47

0.07

6.45

0.03

6.37

0.06

6.37

0.06

6.47

0.07

14

Si

7.51

0.03

7.52

0.06

7.50

0.03

7.51

0.03

7.53

0.03

7.52

0.06

7.52

0.06

7.51

0.04

7.51

0.04

7.51

0.04

7.54

0.05

15

P

5.41

0.03

5.42

0.03

5.41

0.03

5.42

0.03

5.46

0.04

5.46

0.04

5.46

0.04

5.41

0.03

5.36

0.04

5.36

0.04

5.49

0.05

16

S

7.15

0.05

7.12

0.05

7.12

0.03

7.12

0.03

7.16

0.05

7.16

0.05

7.14

0.01

7.12

0.03

7.14

0.05

7.14

0.05

7.21

0.04

17

Cl

[5.25]sp

0.12

[5.25]sp

0.12

5.50

0.30

5.50

0.30

5.50

0.30

5.50

0.30

5.50

0.30

5.50

0.30

18

Ar

[6.50]in

0.10

[6.50]in

0.10

6.50

0.10

6.50

0.10

6.40

0.13

6.18

0.08

6.18

0.08

6.55

0.08

19

K

5.04

0.05

5.10

0.07

5.14

0.03

5.04

0.05

5.06

0.05

5.11

0.09

5.11

0.09

5.12

0.03

5.03

0.09

5.08

0.07

5.08

0.07

5.12

20

Ca

6.32

0.03

6.34

0.06

6.34

0.08

6.32

0.03

6.34

0.03

6.33

0.07

6.33

0.07

6.34

0.04

6.31

0.04

6.31

0.04

6.36

0.02

21

Sc

3.16

0.04

3.15

0.06

3.07

0.04

3.15

0.06

3.16

0.04

3.22

0.04

3.10

0.10

3.10

0.10

3.14

0.04

3.17

0.10

3.05

0.08

3.17

0.10

22

Ti

4.93

0.04

4.97

0.04

4.98

0.05

4.97

0.04

4.93

0.04

4.99

0.04

4.90

0.06

4.90

0.06

4.95

0.05

4.90

0.06

4.90

0.06

5.02

0.06

23

V

3.89

0.08

3.96

0.04

3.96

0.04

3.89

0.08

4.07

0.08

4.00

0.02

4.00

0.02

3.95

0.10

4.00

0.02

4.00

0.02

4.00

0.02

24

Cr

5.62

0.04

5.62

0.07

5.62

0.07

5.62

0.04

5.65

0.04

5.64

0.01

5.64

0.01

5.64

0.04

5.64

0.10

5.64

0.10

5.64

0.13

25

Mn

5.52

0.03

5.45

0.05

5.45

0.05

5.42

0.04

5.47

0.04

5.37

0.05

5.37

0.05

5.43

0.04

5.39

0.03

5.39

0.03

5.39

0.03

26

Fe

7.48

0.04

7.52

0.05

7.52

0.05

7.47

0.04

7.5

0.04

7.48

0.06

7.52

0.06

7.45

0.08

7.50

0.04

7.45

0.05

7.45

0.05

7.45

0.08

27

Co

4.93

0.05

4.96

0.06

4.96

0.06

4.93

0.05

4.99

0.05

4.92

0.08

4.92

0.08

4.99

0.07

4.92

0.08

4.92

0.08

4.92

0.08

28

Ni

6.20

0.04

6.28

0.06

6.28

0.06

6.20

0.04

6.24

0.04

6.23

0.04

6.23

0.04

6.22

0.04

6.23

0.04

6.23

0.04

6.22

0.13

29

Cu

4.18

0.05

4.21

0.03

4.10

0.06

4.21

0.03

4.18

0.05

0.05

4.21

0.04

4.21

0.04

4.19

0.04

4.21

0.04

4.21

0.04

4.21

0.04

30

Zn

4.56

0.05

4.61

0.09

4.61

0.09

4.56

0.05

0.05

4.62

0.15

4.62

0.15

4.56

0.05

4.60

0.03

4.60

0.03

4.62

0.15

31

Ga

3.02

0.05

3.09

0.05

3.02

0.05

0.05

2.88

0.10

2.88

0.10

3.04

0.09

2.88

0.10

2.88

0.10

2.88

0.10

32

Ge

3.63

0.07

3.62

0.07

3.63

0.07

0.07

3.58

0.05

3.58

0.05

3.65

0.10

3.58

0.05

3.58

0.05

3.58

0.05

33

As

34

Se

35

Br

36

Kr

[3.25]in

0.08

[3.25]in

0.08

3.25

0.06

0.06

3.28

0.08

3.28

0.08

3.25

0.06

3.25

0.08

3.28

0.08

3.28

0.08

37

Rb

2.47

0.07

2.57

0.07

2.47

0.07

2.57

0.07

2.60

0.10

2.60

0.10

2.52

0.10

2.60

0.15

2.60

0.15

2.60

38

Sr

2.83

0.06

2.92

0.05

2.95

0.08

2.83

0.06

2.85

0.06

2.92

0.05

2.92

0.05

2.87

0.07

2.92

0.05

2.92

0.05

2.92

0.05

39

Y

2.21

0.05

2.20

0.05

2.21

0.05

2.20

0.05

2.21

0.02

2.21

0.02

2.21

0.05

2.21

0.02

2.21

0.02

2.21

0.02

40

Zr

2.59

0.04

2.59

0.06

2.60

0.10

2.59

0.04

2.59

0.04

2.62

0.06

2.58

0.02

2.58

0.04

2.58

0.02

2.59

0.04

2.59

0.04

41

Nb

1.47

0.06

1.49

0.06

1.47

0.06

1.49

0.06

1.44

0.06

1.42

0.06

1.46

0.04

1.42

0.06

1.42

0.06

1.42

0.06

42

Mo

1.88

0.09

2.04

0.05

1.88

0.09

2.04

0.09

1.92

0.05

1.92

0.05

1.88

0.08

1.92

0.05

1.92

0.05

1.92

0.05

43

Tc

44

Ru

1.75

0.08

1.91

0.10

1.75

0.08

1.91

0.08

1.72

0.10

1.84

0.07

1.75

0.08

1.84

0.07

1.84

0.07

1.84

0.07

45

Rh

0.89

0.08

1.07

0.08

0.89

0.08

1.07

0.08

1.12

0.12

1.12

0.12

0.91

0.10

1.12

0.12

1.12

0.12

1.12

0.12

46

Pd

1.55

0.06

1.61

0.06

1.55

0.06

1.61

0.06

1.66

0.04

1.66

0.04

1.57

0.10

1.66

0.04

1.69

0.04

1.69

0.04

47

Ag

0.96

0.1

1.04

0.10

0.96

0.1

1.04

0.1

0.94

0.30

0.94

0.30

0.94

0.10

0.94

0.25

0.94

0.24

0.94

uncert.

48

Cd

1.77

0.15

1.79

0.15

1.77

0.15

1.79

0.15

1.77

0.11

1.77

0.11

1.77

0.11

1.77

0.11

1.77

0.11

49

In

[0.8]sp

0.20

0.8

0.2

«/ul>1.5

«/ul>1.5

0.80

0.20

1.60

0.20

1.60

0.20

1.56

0.20

50

Sn

2.02

0.1

2.04

0.10

2.02

0.1

2.04

0.1

2.00

0.30

2.00

0.30

2.04

0.10

2.00

0.30

2.00

0.30

2.00

0.30

51

Sb

1.00

0.30

1.00

0.30

1.00

0.30

1.00

0.30

1.00

0.30

52

Te

53

I

54

Xe

[2.25]in

0.08

[2.25]in

0.08

[2.24]

0.06

[2.24]

0.06

2.27

0.08

2.27

0.08

2.24

0.06

2.24

0.02

2.27

0.02

2.27

0.02

55

Cs

56

Ba

2.25

0.07

2.19

0.07

2.20

0.04

2.25

0.07

2.18

0.07

2.20

0.1

2.17

0.07

2.18

0.09

2.17

0.07

2.17

0.07

2.17

0.07

57

La

1.11

0.04

1.14

0.03

1.14

0.03

1.11

0.04

1.14

0.04

1.14

0.03

1.14

0.03

1.10

0.03

1.13

0.05

1.13

0.05

1.13

0.03

58

Ce

1.58

0.04

1.61

0.06

1.61

0.06

1.58

0.04

1.61

0.04

1.61

0.06

1.61

0.06

1.58

0.04

1.70

0.10

1.58

0.09

1.58

0.09

59

Pr

0.72

0.04

0.76

0.04

0.76

0.04

0.72

0.04

0.76

0.04

0.76

0.04

0.76

0.04

0.72

0.04

0.58

0.10

0.71

0.08

0.71

0.08

60

Nd

1.42

0.04

1.45

0.05

1.45

0.05

1.42

0.04

1.45

0.04

1.45

0.05

1.45

0.05

1.42

0.04

1.45

0.05

1.45

0.05

1.50

0.12

61

Pm

62

Sm

0.95

0.04

1.00

0.05

1

0.05

0.95

0.04

0.99

0.04

1.00

0.05

1.00

0.05

0.96

0.04

1.00

0.03

1.01

0.06

0.99

63

Eu

0.52

0.04

0.52

0.04

0.53

0.02

0.52

0.04

0.52

0.04

0.55

0.04

0.52

0.04

0.52

0.03

0.52

0.04

0.52

0.04

0.52

0.06

0.52

0.06

0.52

0.04

64

Gd

1.08

0.04

1.11

0.05

1.11

0.05

1.08

0.04

1.11

0.04

1.11

0.05

1.11

0.05

1.07

0.04

1.11

0.03

1.12

0.04

1.12

0.04

65

Tb

0.31

0.1

0.28

0.10

0.28

0.10

0.31

0.1

0.28

0.1

0.28

0.05

0.28

0.10

0.30

0.10

0.28

0.30

0.28

0.30

0.28

0.30

66

Dy

1.1

0.04

1.13

0.06

1.13

0.06

1.1

0.04

1.13

0.04

1.13

0.06

1.13

0.06

1.10

0.04

1.14

0.08

1.14

0.08

1.14

0.08

67

Ho

0.48

0.11

0.51

0.10

0.51

0.10

0.48

0.11

0.51

0.11

0.51

0.10

0.51

0.10

0.48

0.11

0.51

0.10

0.51

0.10

0.53

0.10

68

Er

0.93

0.05

0.96

0.06

0.96

0.06

0.93

0.05

0.96

0.05

0.96

0.05

0.96

0.06

0.92

0.05

0.93

0.06

0.93

0.06

0.93

0.06

69

Tm

0.11

0.04

0.14

0.04

0.14

0.04

0.11

0.04

0.14

0.04

0.14

0.04

0.14

0.04

0.10

0.04

0.00

0.15

0.00

0.15

0.00

0.15

70

Yb

0.85

0.11

0.86

0.10

0.86

0.10

0.85

0.11

0.86

0.11

0.86

0.10

0.86

0.10

0.84

0.11

1.08

0.15

1.08

0.15

1.08

0.15

71

Lu

0.1

0.09

0.12

0.08

0.12

0.08

0.1

0.09

0.12

0.09

0.12

0.08

0.12

0.08

0.10

0.09

0.06

0.10

0.06

0.10

0.06

0.10

72

Hf

0.85

0.05

0.88

0.03

0.88

0.03

0.85

0.05

0.88

0.05

0.88

0.05

0.87

0.04

0.88

0.08

0.85

0.04

0.88

0.08

0.88

0.08

0.88

0.08

73

Ta

74

W

0.83

0.11

1.03

0.11

0.83

0.11

1.03

0.11

1.11

0.15

1.11

0.15

0.85

0.12

1.11

0.15

1.11

0.15

1.11

0.15

75

Re

76

Os

1.4

0.05

1.5

0.05

1.4

0.05

1.5

0.05

1.36

0.19

1.36

0.19

1.45

0.11

1.40

0.08

1.25

0.11

1.45

0.10

1.45

0.10

77

Ir

1.42

0.07

1.46

0.07

1.42

0.07

1.46

0.07

1.38

0.05

1.38

0.05

1.38

0.07

1.38

0.05

1.38

0.05

1.38

0.05

78

Pt

1.74

0.30

1.74

0.30

1.74

uncert.

79

Au

0.91

0.08

0.93

0.08

0.91

0.08

0.93

0.08

1.01

0.18

1.01

0.18

0.92

0.10

1.01

0.15

1.01

0.15

1.01

0.18

80

Hg

81

Tl

[0.95]sp

[0.2]

[0.9]

0.2

0.95

0.20

0.95

0.20

0.90

0.20

0.90

0.20

0.90

0.20

0.72–1.1

uncert.

82

Pb

1.92

0.08

2.05

0.08

1.92

0.08

2.05

0.08

2.00

0.06

2.00

0.06

1.75

0.10

2.00

0.06

2.00

0.06

2.00

0.06

83

Bi

90

Th

«/ul>0.03

uncert.

«/ul> 0.07

uncert.

0.03

0.1

0.07

0.1

< 0.08

0.08

0.03

0.08

0.03

0.02

0.10

uncert.

uncert.

0.12

uncert.

92

U

< −0.47

< −0.47

< −0.47

< −0.47

< −0.47

Table 4. Elemental Abundances in the Solar Photosphere. Atomic abundances scaled to: 12 + log10(n(E)/n(H)) Part 2.

Z

E

[PJ03]

[GS02]

[H01]

[GS98]

[GNS96]

[GN93]

[AG89]

[G84]

[P79]

[E77]

[RA76]

[W71]

[M68]

[GBB68]

[M65]

[GMA60]

[C51]

[U48]

[U47]

[U45]

[Me43]

[R29]

[P25]

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

A(E)

1

H

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

12

2

He

10.99

10.99

10.99

10.99

10.93

10.99

11

10.8

10.9

10.8

11.2

11

10.5

9.3

3

Li

1.10

1.10

1.10

1.16

1.15

1.16

1

1

1

1

0.6

<0.9

1

1.54

0.96

1.08

2.5

1.0

4

Be

1.40

1.40

1.40

1.15

1.15

1.15

1.15

1.2

1.1

1.15

1.06

2.34

2.33/1.24

2.34

2.36

2.3

5

B

2.55

2.70

2.55

2.60

2.60

2.60

2.60

2.4

2.3

<2.1

<2.8

<3.6

6

C

8.59

8.52

8.592

8.52

8.55

8.55

8.56

8.69

8.66

8.7

8.62

8.57

8.51

8.55

8.62

8.72

7.78

8.64

8.4

7.56

7.9

5.5

7

N

7.93

7.95

7.931

7.92

7.97

7.97

8.05

7.99

7.93

7.9

7.94

8.06

8.06

8.93

8.86

7.98

8.1

8.86

8.73

8.09

8.1

8

O

8.69

8.73

8.736

8.83

8.87

8.87

8.93

8.91

8.88

8.8

8.84

8.83

8.83

8.93

8.86

8.96

8.65

8.22

8.98

8.99

8.56

9.5

7.1

9

F

4.56

4.56

4.56

4.56

4.56

4.56

4.56

4.6

4.6

4.56

4.56

10

Ne

8.00

8.06

8.001

8.08

8.08

8.08

8.09

8.00

7.8

7.7

7.57

7.45

9.06

9.06

11

Na

6.33

6.33

6.33

6.33

6.33

6.33

6.33

6.29

6.3

6.28

6.24

6.3

6.30

6.33

5.77

6.53

6.36

6.56

7.7

6.2

12

Mg

7.54

7.58

7.538

7.58

7.58

7.58

7.58

7.58

7.6

7.6

7.6

7.54

7.36

7.40

7.57

7

7.76

7.79

8.39

8.3

6.6

13

Al

6.47

6.47

6.47

6.47

6.47

6.47

6.47

6.4

6.4

6.52

6.4

6.2

6.20

6.17

5.82

6.58

6.58

6.39

6.9

6.0

14

Si

7.54

7.55

7.536

7.55

7.55

7.55

7.55

7.55

7.6

7.6

7.65

7.55

7.70/7.24

7.74

7.5

7.50

7.12

6.78

7.7

7.75

7.87

7.8

6.6

15

P

5.45

5.43

5.45

5.45

5.45

5.45

5.45

5.5

5.5

5.5

5.43

5.34

5.34

16

S

7.33

7.33

7.33

7.33

7.21

7.21

7.21

7.2

7.2

7.2

7.21

7.3

7.30

6.41

7.17

7.17

7.57

6.2

17

Cl

5.5

5.5

5.5

5.5

5.5

5.5

5.5

5.4

5.5

5.5

5.65

18

Ar

6.4

6.4

6.4

6.52

6.52

6.56

6.58

6.4

6.0

6.00

6.65

19

K

5.12

5.12

5.12

5.12

5.12

5.12

5.12

5.15

5.2

5.16

5.05

4.7

4.70

5.01

4.69

5.45

5.09

7.3

4.4

20

Ca

6.36

6.36

6.36

6.36

6.36

6.36

6.36

6.3

6.3

6.35

6.33

6.04

6.15

6.46

5.72

6.48

6.48

6.57

7.2

5.8

21

Sc

3.17

3.17

3.17

3.17

3.17

3.1

3.1

3.0

3.1

3.04

3.07

2.85

3.24

2.8

2.82

2.82

3.58

4.1

22

Ti

5.02

5.02

5.02

5.02

5.02

4.99

5.02

4.8

5.0

5.05

4.74

4.63/4.81

4.85

4.58

4.68

7.56

4.45

5.21

4.57

5.7

5.1

23

V

4.00

4.00

4.00

4.00

4.00

4.00

4.00

4.0

4.1

4.02

4.1

4.17

4.28

4.12

3.70

3.54

4.3

4.09

5.5

4.0

24

Cr

5.67

5.67

5.67

5.67

5.67

5.67

5.67

5.7

5.7

5.71

5.7

5.12/5.01

5.12

5.07

5.36

5.07

5.83

4.87

6.2

4.9

25

Mn

5.39

5.39

5.39

5.39

5.39

5.39

5.45

5.4

5.4

5.42

5.2

4.85

5.01

4.80

4.90

4.95

5.71

5.09

6.4

5.6

26

Fe

7.45

7.5

7.448

7.50

7.50

7.50

7.67

7.67

7.5

7.6

7.50

7.40

6.80

6.79

6.70

6.57

7.16

7.21

7.97

7.97

6.99

7.7

5.8

27

Co

4.92

4.92

4.92

4.92

4.92

4.92

4.92

5.0

5.0

4.90

4.50

4.70

4.56

4.40

4.64

4.52

5.28

4.69

6.1

28

Ni

6.25

6.25

6.25

6.25

6.25

6.25

6.25

6.3

6.3

6.28

6.28

5.77

5.64

5.44

5.91

3.64

6.2

6.39

6.5

29

Cu

4.21

4.21

4.21

4.21

4.21

4.21

4.21

4.1

4.1

4.06

4.45

4.45

4.18

3.5

5.04

4.8

3.72

4.48

4.39

5.5

30

Zn

4.6

4.6

4.6

4.6

4.60

4.6

4.6

4.4

4.4

4.45

4.42

3.52

3.23

3.52

4.40

4.52

4.27

5.03

5.57

5.4

5.2

31

Ga

2.88

2.88

2.88

2.88

2.88

2.88

2.88

2.8

2.8

2.8

2.84

2.72

2.36

2.5

32

Ge

3.41

3.41

3.41

3.41

3.41

3.41

3.63

3.4

3.4

3.5

3.32

2.49

3.29

3.5

33

As

34

Se

35

Br

36

Kr

3.3

3.31

3.31

37

Rb

2.6

2.6

2.6

2.60

2.60

2.60

2.60

2.6

2.6

2.6

2.63

2.48

2.48

2.2

38

Sr

2.97

2.92

2.97

2.97

2.90

2.90

2.90

2.9

2.9

2.9

2.82

3.02

2.60

2.88

2.84

3.7

3.8

2.6

39

Y

2.21

2.24

2.24

2.24

2.24

2.24

2.24

2.1

2.1

2.1

1.62

3.2

2.33/2.42

3.2

2.25

2.7

3.56

3.1

40

Zr

2.59

2.6

2.6

2.60

2.60

2.60

2.56

2.8

2.8

2.75

2.65

3.31/2.28

2.65

2.23

1.86

2.72

3

41

Nb

1.42

1.42

1.42

1.42

1.42

1.42

2.1

2.0

2.0

1.9

2.3

2.3

2.73

2.3

1.95

1.5

42

Mo

1.92

1.92

1.92

1.92

1.92

1.92

1.92

2.2

2.2

2.16

1.9

2.3

2.52

2.3

1.90

1.27

2.13

1.9

43

Tc

44

Ru

1.84

1.84

1.84

1.84

1.84

1.84

1.84

1.8

1.9

1.83

1.57

1.82

1.98

1.82

1.43

2.2

45

Rh

1.12

1.12

1.12

1.12

1.12

1.12

1.12

1.5

1.5

1.4

1.55

1.37

1.55

1.37

0.78

1

46

Pd

1.69

1.69

1.69

1.69

1.69

1.69

1.69

1.5

1,5

1.5

1.57

1.57

1.64

1.27

1.21

1.6

47

Ag

0.94

0.94

0.94

0.94

0.94

0.94

0.94

0.9

0.9

0.85

0.67

0.75

1.12

1.04

0.14

1.5

48

Cd

1.77

1.77

1.77

1.77

1.77

1.86

1.86

2

2

1.85

1.97

1.54

2.64

1.66

1.46

2.7

49

In

1.66

1.66

1.66

1.66

1.66

1.66

1.66

1.7

1.7

1.65

1.71

1.45

1.68

1.28

1.16

0.5

50

Sn

2.00

2.00

2.00

2.00

2.00

2.00

2.00

2

2

2

1.71

1.54/2.05

1.32

2.05

1.54

1.7

51

Sb

1.00

1.00

1.00

1.00

1.00

1.00

1.00

1

1

1

0.75

1.94/.42

1.04

0.42

1.94

1.3

52

Te

53

I

54

Xe

2.16

2.17

2.17

55

Cs

1.12

<1.9

<2.1

<1.9

<1.79

56

Ba

2.13

2.13

2.13

2.13

2.13

2.13

2.13

2.1

2.1

2.09

2.22

2.1

2.06

2.5

2.10

2.38

2.44

3.8

2.1

57

La

1.17

1.17

1.17

1.17

1.22

1.22

1.22

1.1

1.1

1.13

1.11

2.03

2.03

2.3

58

Ce

1.58

1.58

1.58

1.58

1.55

1.55

1.55

1.6

1.6

1.55

1.62

1.78

1.98

2.9

59

Pr

0.71

0.71

0.71

0.71

0.71

0.71

0.71

0.8

0.8

0.66

0.78

1.45

1.59

1.1

60

Nd

1.50

1.50

1.50

1.50

1.50

1.50

1.34

1.2

1.2

1.23

1.44

1.93

2.5

61

Pm

62

Sm

1.01

1.01

1.01

1.01

1.01

1.00

0.80

0.7

0.7

0.72

0.91

1.62

2

63

Eu

0.51

0.51

0.51

0.51

0.51

0.51

0.51

0.7

0.7

0.7

0.51

0.96

1.9

64

Gd

1.12

1.12

1.12

1.12

1.12

1.12

1.12

1.1

1.1

1.12

1.15

1.13

1.6

65

Tb

−0.1

−0.1

−0.1

−0.1

−0.1

−0.1

0.2

0.42

66

Dy

1.14

1.14

1.14

1.14

1.14

1.10

1.10

1.1

1.1

1.06

1.11

1

2.1

67

Ho

0.26

0.26

0.26

0.26

0.26

0.26

0.26

0.5

68

Er

0.93

0.93

0.93

0.93

0.93

0.93

0.93

0.8

0.8

0.76

0.89

0.6

69

Tm

0.00

0.00

0.00

0.00

0.00

0.00

0.00

0.3

0.3

0.26

0.09

70

Yb

1.08

1.08

1.08

1.08

1.08

1.08

1.08

0.8

0.2

0.9

0.87

1.53

2.33(I), 0.49(II)

71

Lu

0.06

0.06

0.06

0.76

0.76

0.76

0.76

0.8

0.8

0.76

0.09

72

Hf

0.88

0.88

0.88

0.88

0.88

0.88

0.88

0.9

0.9

0.8

0.9

73

Ta

74

W

1.11

1.11

1.11

1.11

1.11

1.11

1.11

0.8

0.8

1.7

2.57

0.7

75

Re

<−0.3

<−0.3

<−0.3

0.25

76

Os

1.45

1.45

1.45

1.45

1.45

1.45

1.45

0.7

0.7

0.7

1.34

0.75

77

Ir

1.35

1.35

1.35

1.35

1.35

1.35

1.35

2.2

0.8

0.85

1.24

2.21

0.3

78

Pt

1.80

1.80

1.80

1.80

1.80

1.80

1.80

1.8

1.8

1.75

1.75

2.1

79

Au

1.01

1.01

1.01

1.01

1.01

1.01

1.13

0.1

0.8

0.75

0.91

0.8

80

Hg

<2.1

<2.1

<2.1

2.61

1.13?

81

Tl

0.90

0.90

0.90

0.90

0.90

0.90

0.90

0.9

0.9

0.9

0.81

1.9

82

Pb

2.00

1.95

1.95

1.95

1.95

1.85

1.9

1.9

1.9

1.93

1.75

1.63

1.66

1.63

1.33

2.04

1.7

83

Bi

<1.9

<1.9

<1.9

0.78

0.81

90

Th

0.12

0.12

0.02

0.2

0.2

0.2

−0.02

92

U

<−0.47

<−0.47

<−0.47

<−0.47

<−0.47

<−0.47

<−0.47

<0.6

<0.6

<0.6

−0.5

a Notes: Atomic abundances scaled to: A(E)=12+log10(n(E)/n(H)). Abundances from photospheric data unless indicated: “in” = indirect determination. “sp” sunspot. Helium from helioseismology; see text. Mass fractions for recommended 3D composition: X=0.7389 (Basu & Antia, 2004), Y=0.2462 for Z=0.0.0149 and Z/X=0.0201. Mass fractions for recommended 1D composition: X=0.7389 (Basu & Antia, 2004), Y=0.2458 for Z=0.0153 and Z/X=0.0207. Neon estimated from solar wind and coronal Ne/O ratios from Young (2018) using oxygen abundances recommended here, Ne abundance is still uncertain and debatable. Argon from solar wind, B-stars, and other considerations, see Lodders (2008). Krypton and Xenon estimated from nuclear systematics; see text. Recommended 1D: largely from SL+, SSG15, Z16. Recommended 3D largely from C11 and SSG15, plus Amarsi, Barklem, Asplund, Collet, and Zatarinny (2018) (O), Caffau et al. (2015) (O), Amarsi and Asplund (2017) (Si), Armarsi et al. (2019) (C), Caffau et al. (2011) (C), Bergemann et al. (2017) (Mg), Bergemann et al. (in press) (Mn), Lind et al. (2017) (Fe), Nordlander et al. (2017) (Al). Recommended 1D largely from H01, SSG15, Z16, Alexeeva and Mashonkina (2015) (C). Sunspot data: Maas, Pilachowski, and Hinkle (2016) (Cl), Maiorca et al. (2014) (F), Vitas et al. (2008) (In). Reference codes in the table: [AG89] Anders and Grevesse (1989); [AGS05, AGSS09] Asplund et al. (2005, 2009); [C11] Caffau et al. (2011); [C51] Claas (1951); [E77] Envold (1977); [G84] Grevesse (1984); [GAS07], Grevesse, Asplund, and Sauval (2007); [GBB68] Grevesse, Blanquest, & Boury (1968); [GMA60] Goldberg et al. (1960); [GNS96] Grevesse, Noels, & Sauval (1996); [GS02] Grevesse and Sauval (2002); [GS93] Grevesse and Noels (1993); [GS98] Grevesse and Sauval (1998); [H01] Holweger (2001); [L03] Lodders (2003); [LPG09] Lodders et al. (2009); [M65, M68] Müller (1965, 1968); [Me43] Goldberg & Aller, 1943, p. 114; [P76], Pagel (1979); [PJ03] Palme and Jones (2003); [PLJ14] Palme et al. (2014); [P25] Payne (1925a, 1925b); stellar abundances; [R29] Russell (1929); [RA76] Ross and Aller (1976); [SSG15] Scott, Asplund, et al. (2015); Scott, Grevesse, et al. (2015); Grevesse et al. (2015); [SL+] Sneden et al. (2016); Sneden et al. (2009); Lawler et al. (2014, 2017, 2018, 2019); [U45, U47] Unsöld (1945, 1947) (as quoted in Suess, 1947b). [U48] Unsöld (1948); [W71] Withbroe (1971); and [Z16] Zhao et al. (2016).

Coronal Abundances

The chromosphere is the transition zone between the photosphere (with a brightness temperature of around 5800K) and the tenuous outer corona (with a kinetic gas temperature around 2 million K). The chromosphere has a large temperature gradient and is heterogeneous. Chromospheric emission spectra obtained during solar eclipses show the presence of neutral and singly ionized atoms, whereas quiet coronal emissions show the forbidden lines of highly ionized atoms in the visible regions of the spectra and permitted lines in the extreme ultraviolet (EUV) and X-ray wavelengths. Analyses of chromospheric plasma spectra require additional knowledge for atomic properties, including ionization equilibria and deviations from them, electron densities, and line excitation mechanisms. Elemental abundances for the corona are presented in Table 5.

Coronal abundances are not easy to quantify because atomic data for radiative transition probabilities and cross-sections for excitation, ionization, and recombination processes must be available from laboratory measurements. First, coronal abundances were derived from observations of forbidden lines and later from extreme UV observations. The first spectroscopic analyses for 20 elements in the solar corona was achieved by Menzel (1931); later abundances for the corona from optical observations are by, for example, Woolley and Allen (1948) and Pottasch (1964b); see Feldman and Widing (2003) for more studies. In 1946 the U.S. Naval Research Laboratory mounted a spectrograph on the first V-2 rocket launched by the United States to obtain the first EUV spectra of the Sun (see Trousey [1967] for a review of early solar UV research). Subsequently Pottasch (1964a, 1964c, 1967) analyzed solar EUV spectra and found iron was 40 times higher in the corona than in the photosphere but closer to meteoritic values, which was one of the indications of the erroneous photospheric iron abundance (see previous discussion).

EUV and X-ray spectroscopy with instruments on board of several spacecraft recently provided measurements of coronal matter in solar flares, for example, the SOHO Coronal Diagnostic Spectrometer (Mason et al., 1999), the Hinode EUV Imaging Spectrometer (O’Dwyer et al., 2011), the solar X-ray spectra from the REntgenovsky Spektrometr s Izognutymi Kristalami crystal spectrometer on the CORONAS-F spacecraft (Sylwester et al., 2015), and Solar Assembly for X-rays on the Mercury MESSENGER spacecraft (Dennis et al., 2015). For EUV and X-ray spectroscopic observations of the Sun by an armada of spacecraft, see Milligan and Ireland (2018).

Solar Wind Abundances

The solar wind stems from mass ejections originating from the photosphere and the corona. The solar wind is an interplanetary presence of solar material, and the term heliosphere is used to describe where solar wind and the solar magnetic field dominate over the interstellar particle flow (i.e., galactic cosmic rays) and the interstellar magnetic field. The solar wind as corpuscular irradiation was first recognized in order to explain why cometary tails always flow outwards with motion perpendicular from the Sun’s sphere (Biermann, 1951; Parker, 1958). Implanted solar wind was first detected in noble gas rich meteorites, later in lunar rocks and soil that had been exposed to the solar wind, and directly captured in the foils posted on the lunar surface during the Apollo missions (Geiss et al., 2004). There are limitations to obtaining abundances for many elements from surface-exposed rocks due to the ubiquitous presence of the same elements in the lunar and meteoritic rocks themselves. Useful elemental and isotope data for solar wind can only be obtained for the noble gases and nitrogen, which are typically not part of indigenous minerals in meteoritic and lunar rocks (see Pillinger [1979] for a review). The foil experiments may not suffer from the problems of ascribing an element occurrence to the solar wind; however, the exposure times of the foils were relatively short and only major abundant elements could be collected. Space missions provided in-situ measurements of the solar wind compositions in the heliosphere with energy/charge spectrometers; these include the Advanced Composition Explorer and Ulysses (see, e.g., Gloeckler et al., 1998; Gloeckler & Geiss, 2007). Elemental abundances from these experiments are given in Table 5. The most ambitious mission with sample return and laboratory analyses of solar wind samples is the Genesis mission (Burnett & Genesis Science Team, 2011; Burnett, Jurewicz, & Woolum, 2019). The analyses have already provided well-characterized noble gas abundances (e.g., Crowther & Gilmour, 2013; Heber et al., 2009; Heber et al., 2012; Meshik, Hohenberg, Pravdivtseva, & Burnett, 2014; Pepin, Schlutter, Becker, & Reisenfeld, 2012; Vogel et al., 2011) and C, N, and O abundances (e.g., Heber et al., 2013; Laming et al., 2017), whereas the analyses of rock-forming elements (e.g., Burnett et al., 2017; Burnett et al., 2019; Heber et al., 2014; Jurewicz et al., 2011; Jurewicz, Olinger, Burnett, Rieck, & Woolum, 2019; Rieck et al., 2016; Westphal, Huss, Koeman-Shields, Welten, & Westphal, 2019) are ongoing. These preliminary abundances are listed in Table 5.

Figures 3a and 3b show coronal and solar wind abundances. The solar wind measurements pertain to ejected solar matter. The solar wind consists of a plasma with highly ionized atoms. The ionization state changes as plasma is heated and accelerated and slow and fast plasma streamers and coronal mass ejections develop. The slow and fast solar wind and SEPs are related to these phenomena. The elemental and isotopic abundances in the corona and solar wind are fractionated from photospheric abundances due to hydromagnetic processes that depend on the charge state or charge-to-mass ratios of individual atoms.

The fractionation occurs when plasma flows from the photosphere to the chromosphere into the heliosphere and appears to be a function of the first ionization potential (FIP) of the elements. As shown in Figure 3a, elements with low FIP (below 10 eV, including Na, K, Mg, Al, Si, Fe) have more or less uniformly higher abundances in coronal and solar wind matter than observed in the photosphere, in contrast to elements with high FIP above 10eV, including C, N, O, and noble gases (e.g., Reisenfeld et al., 2013; von Steiger & Zurbuchen, 2016). Fast solar wind originates from coronal holes and is closer to photospheric abundances, but elements with low and high FIP are still fractionated, although not as much as the slow solar wind, which makes it difficult to use these solar wind abundances to derive a representative solar composition.

More energetic forms (> 1 MeV/nucleon) of the solar wind present within the interplanetary medium of the heliosphere are so-called SEPs, which occur as either “gradual” or “impulsive” types (Desai & Giacalone, 2016). The gradual SEPs are created when the slow solar wind is accelerated by shocks from coronal mass ejections as matter passes through the corona. These SEP abundances are sometimes taken as representative for the corona. In contrast, impulsive solar flares or impulsive SEP events resulting from flare events, are compositionally more variable, and could be regarded as extreme forms of the solar wind. Elemental composition for the SEPs have been measured, for example, with the cosmic ray subsystem aboard the Voyager 1 and 2 spacecraft (see Desai & Giacalone, 2016; Reames, 2018). Interplanetary coronal mass ejections are highly energetic events, and solar material in them is highly fractionated from photospheric abundances. Zurbuchen et al. (2016) analyzed data from the Solar Wind Ion Composition Spectrometer on board the Advanced Composition Explorer (ACE). Elemental abundances for these solar wind components are included in Table 5.

Figure 3a. Elemental abundances in the model corona (Schmelz et al., 2012). The solar wind as measured by the Genesis mission (see Table 5 for detailed references) and in SEPs (Reames, 2018) normalized to Mg and photospheric abundances (3D from Table 4), as a function of the FIP of the elements. Error bars include the uncertainties from the photospheric element/Mg ratios used for normalization. See Figure 3b for the same data plotted as a function of FIT and text for explanations.

Figure 3b. Same as Figure 3a, but normalized abundances are plotted versus the FIT. The FIT data are from Marsch et al. (1995), Geiss (1998), and Giammanco et al. (2008); note that FIT values for Mg and Fe were estimated based on FIP.

Enrichments of elements with low FIP may arise from Alfenic waves causing ponderomotive forces (forces in spatially inhomogeneous, electromagnetic fields acting on moving charged particles) that accelerate chromospheric ions (see, e.g., Dahlburg, Laming, Taylor, & Obenschain, 2016; Laming, 2015; Reames, 2018). However, the compositional differences between the fast and slow solar wind are not yet fully understood with existing models on the origins of the various solar wind regimes. In addition to the dependence on FIP, there is an apparent correlation of the elemental enrichments with first ionization time (FIT; Geiss & Bochsler, 1986; Widing & Feldman, 2001).

Elements with low FIP are two to four times more abundant in the corona than in the photosphere, whereas elements with high FIP have comparable abundances when the ratio corona to photosphere is also normalized to a high FIP element such as H or O. Conversely, when the abundance ratio is normalized to a low FIP element (such as Mg as done in Figures 3a and 3b), all low FIP elements have close to photospheric abundances around unity, whereas the high FIP elements are below unity, indicating an apparent depletion. The “hybrid” abundances for the corona from Schmelz, Reames, von Steiger, and Basu (2012) are modeled by assuming low FIP elements (Mg, Fe, etc.) have essentially constant element/hydrogen ratios of 2.14 times the photospheric ratios and high FIP elements (C, N, O, noble gases) ratios of 0.71, respectively. Elements with intermediate FIP (P, S, Zn) are in between. This approach places hydrogen, a high FIP element, at unity (the photospheric value), above all other high FIP elements (which are at 0.71 times the photospheric value) by Schmelz et al. This is peculiar because H and O have similar FIPs, and therefore H and O should show similar abundances relative to the photosphere if the FIP is the underlying cause for the element fractionations between the photosphere and the corona. The coronal abundances by Feldman and Widing (2003) place the low FIP elements at 2.1 times the photospheric abundance ratios and the high FIP elements, including hydrogen, at the photospheric level.

The FIP may not be the main governing parameter, and the FIT may be more important. The FIT describes the time after which 50% of an element has reached a certain ionization stage and the effectiveness with which atoms can be fractionated in the competition between photoionization and dynamic atom-ion processes. It is not a fixed atomic property but a model-dependent and time-variable convolution of the ionization cross-section with the solar EUV spectrum (e.g., Geiss, 1998; Geiss, Gloeckler, & von Steiger, 1994; von Steiger & Schwadron, 2000). For example, the FIT might be a better ordering parameter for elements with similar FIP, if they show abundance fractionations. The same data from Figure 3a are plotted as a function of FIT in Figure 3b (see also Vogel et al. [2011] for similar figures for the solar wind). Overall, the two figures show similar trends, and, qualitatively, the correlation with FIP is more like a step-function than a steady correlation with FIT. The correlation with FIT shows more scatter, which is in part due to the difficulty in deriving estimates for the FIT, which can vary greatly (see Geiss, 1998; Giammanco, Wurz, & Karrer, 2008; Marsch, von Steiger, & Bochsler, 1995). The uncertainties in some abundance measurements remain relatively large for various solar wind regimes (see e.g., similar diagrams in Vogel et al., 2011) but also for photospheric values used for normalization (Table 4). Overall, the existing and anticipated high-quality results of the Genesis mission should help to clarify the issue whether FIP, FIT, or yet another mechanism is a key parameter of the fractionation of the solar wind from photospheric material.

Table 5. Elemental Abundances in the Solar Corona and the Solar Wind Normalized to Mg = 100 Atoms

Z

FIP, eV

Photo-sphere

Corona (model)

Corona

Corona

Corona (optical)

Corona (EUV)

Corona (X-ray)

Gradual SEP

Gradual SEP

3He-rich SEP

Slow Solar Wind

Fast Solar Wind

Coronal Mass Ejections

Solar Wind

Solar Wind

(Table 4)

[Sch12]

[M85ab]

[W71]

[FW03]

[W75]

[VP81]

[R18]

[M85ab]

[Mas04]

[SSWave]

[FSWave]

[CMEave]

[SWBulk]

[Genesis]

2

H

13.598

2.75E+6

1.35E+6

2.68E+6

2.24E+6

*

3.16E+6

2.24E+6

1.00E+6

1.68E+6

1.81E+6

9.76E+5

1.51E+6

1.02E+6

2

He

24.587

2.34E+5

8.05E+4

2.63E+5

*

5.11E+4

3.09E+4

1.46E+4

(6.06±0.03)E+4

(7.10±0.03)E+4

(4.79±0.09)E+4

5.94E+4

(4.87±0.27)E+4

6

C

11.260

813

303

632

1260

*

1200

236

236

87

450±49

635±27

303±7

428

383±19

7

N

14.534

195

69

105

251

*

141

72

66

35

71±25

126±13

74±4

103

69±5

8

O

13.618

1413

514

663

1259

*

660

631

562

528

270

693±49

918±50

531±22

687

723±47

10

Ne

21.565

389

107

95

79

*

51

100

88

68

70

77±13

96±34

88

75

86±5

11

Na

5.139

4.5

5.8

7.4

6.3

5.5

5.8

7.1

5.8

6.9

6.1

4.8

6.1±1.2

12

Mg

7.646

100

100

100

100

100

100

100

100

100

100

100

100

100

100

100±6

13

Al

5.986

7.4

8.5

7.4

6.3

8.1

5.6

8.8

7.2

8.1

7.6

14

Si

8.152

89

95

105

100

95

100

105

85

81

111

109±13

104±27

94±3

118

15

P

10.487

0.71

0.4

0.60

*

0.8

0.37

1

0.95

16

S

10.360

39

23

23

32

*

35

18

14

16

32

39±6

41±11

35±2

38

18

Ar

15.760

8.7

3.0

5.7

13

*

17

5.4

2.4

3.1

2.1

2.3±0.8

1.90±0.09

19

K

4.341

0.33

0.37

1.3

*

0.31

0.31±0.05

20

Ca

6.113

5.8

6.2

7.9

5.6

5.9

3

7.2

6.2

6.2

16

5.5

5.7±1

24

Cr

6.767

1.15

1.26

1.60

*

1.2

1.8

1.4

1.4

25

Mn

7.434

0.91

0.7

0.80

*

0.89±0.10

26

Fe

7.902

83

87

105

105

83

85

100

74

80

260

90±15

73±10

85±9

88

85±6

28

Ni

7.640

4.40

4.9

5.8

7.6

4.7

6.3

5.6

3.6

3.7

4.4

36

Kr

14.000

4.57E-3

1.82E-3

(1.24±0.10)E-3

54

Xe

12.130

4.90E-4

1.78E-4

(2.76±0.15)E-4

Note: [FW03] Feldman and Widing (2003), optical spectra/forbidden lines. Entries marked * indicate that photospheric abundances were assumed in their table. [M85ab] Meyer (1985a, 1985b) as recommended by the author. [Mas04] Mason et al. (2004). 3He-rich SEP, impulsive flares. [R18] Reames (2018), gradual SEP. [Sch12] Schmelz et al. (2012). Calculated from the adopted photospheric composition here and their relative enrichments/depletion factors, see text. [VP81] Veck and Parkinson (1981), OSO-8 X-ray spectrometer. [W71] Withbroe (1971) as recommended by the author. [W75] Withbroe (1975), EUV spectra. [CMEave] Coronal mass ejections from ACE spacecraft measurements; uncertainty given here probably mainly reflects real temporal variations; Reisenfeld et al. (2013); Zurbuchen et al. (2016). [FSWave] Fast solar wind average (originating from coronal holes) from ACE and Ulysses spacecraft measurements; uncertainty given here probably mainly reflects real temporal variations; Bochsler (2007, 2009); Gloeckler and Geiss (2007); von Steiger et al. (2000); von Steiger and Zurbuchen (2016); Reisenfeld et al. (2013). [SSWave] Slow solar wind (interstream) average from ACE and Ulysses spacecraft measurements; uncertainty given probably mainly reflects real temporal variations; Bochsler (2007, 2009); von Steiger et al. (2000); von Steiger and Zurbuchen (2016); Reisenfeld et al. (2013); Pilleri et al. (2015). [SWBulk] Bulk solar wind, Pilleri et al. (2015); Reisenfeld et al. (2013); see previous comment. [Genesis] Burnett et al. (2017); Rieck et al. (2019) (Mg); Heber et al. (2013) (C,N,O); Jurewicz et al. (2011, 2019) (Fe, Mg); Koeman-Shields et al. (2016) (H); Crowther and Gilmour (2013); Heber et al. (2009); Meshik et al. (2014); Vogel et al. (2011) (Ar, Kr, Xe); Pepin et al. (2012, 2012) (Ne, Ar); Rieck et al. (2016) (Na, K); Reisenfeld et al. (2013) (H, He); Westphal et al. (2019) (Mn).

Other Sources for Solar System Abundances

For some elements, obtaining their abundances by measurements of meteorites or solar environs remain elusive. The noble gases are highly depleted in meteorites, and geochemical analyses from meteorites cannot provide their abundances.

Solar measurements are problematic because spectral lines are absent, weak, or very difficult to model, and physical fractionations as noted earlier for the solar wind make abundance data difficult to interpret without any baseline for reference. For helium, results from helioseismology are used as described earlier.

For neon and argon, spectroscopic observations of the much hotter B-type dwarf stars can provide abundances. These stars have similar relative abundances for elements observable in the solar photosphere, and, therefore, the Ne and Ar abundances can be derived by analogy. However, B-stars are younger than the Sun, and they could have formed from a different mix of elements because of continuing nucleosynthesis in stars and element release into the interstellar medium. Galactic chemical evolution models consider yields of nucleosynthesis models, and, for elements with similar production mechanism, the element abundance ratios should vary in step (just like elements with similar chemical properties fractionate in a similar manner during chemical processing so that element abundance ratios remain about constant).

Another approach to estimate elemental abundances is to consider elemental and isotopic variations in the solar abundance distributions. This type of work is directly related to understanding nuclear structure and nucleosynthetic mechanisms, which was one of the reasons researchers began to look at elemental abundances to begin with. Starting with Oddo’s (1914) and Harkins’ (1917) discovery that elements with odd atomic numbers are less abundant than their even-numbered neighbors, better abundance analyses eventually revealed well-defined structures in elemental abundances, and later in the isotopic abundance distribution curves (e.g., Goldschmidt, 1937). Figure 4, for which the data are described in the following section, shows the abundance curve of the elements as a function of atomic number and illustrates Harkin’s discovery. This and, more importantly, the isotopic composition of the element and isotopic (or nuclide) distributions led to the recognition of nuclide abundance rules by Suess (1947a, 1947b), which then enabled him to predict nuclide abundances and elemental abundances. As already done by Goldschmidt (1937), the abundances of Kr and Xe still rely on interpolations from abundances of their neighboring elements and nuclides, but they are based on much better understood underlying physics about nuclear structure and nucleosynthesis (see, e.g., Anders & Grevesse, 1989; Cameron, 1982; Lodders, 2003; Palme & Beer, 1993). Earlier abundance compilations often used graphical interpolations between neighboring elements (or isotopes) to obtain an abundance estimate for “missing” elements (or isotopes) from plots of elemental abundances versus atomic number (such as Figure 4) or isotope abundances as a function of mass number (such as Figure 5). When employing this method, the best approach is to only use the abundances for isotopes with odd mass numbers because there is usually only one odd-numbered isotope per mass number, and the abundance curve for odd-numbered isotopes is particularly smooth, as noted long ago (Suess, 1947a, 1947b). If the isotopic composition of an element is known, and the abundance of one (or more) of its isotopes can be estimated from the abundance curve, the total elemental abundance is obtained. The same principle, based on better underlying physics, uses nucleosynthesis network calculations to predict the relative contributions from the different nucleosynthetic production process operating in stellar sources. Most successful abundance estimates come from models for the slow-neutron capture process (“s-process”) that operates in evolved low to intermediate mass stars (2–5 solar masses) and produces many isotopes of the elements heavier than iron (e.g., Arlandini et al., 1999; Bisterzo, Travaglio, Wiescher, Käppeler, & Gallino, 2017; Sneden, Cowan, & Gallino, 2008). If an element has one or more isotopes only or largely produced by the s-process (which is the case for, e.g., 82Kr and 128,130Xe), the calculated isotopic abundances from the s-process network can serve to obtain the total elemental abundance, again if the overall isotopic composition of the element is known. Because there are only elements heavier than Fe in the s-process, the anchoring to the solar composition scale relative to 106 silicon atoms is usually done via the s-process-only isotope 150Sm, because its abundance is independently known on the solar abundance scale. Renormalizing the computed s-process abundances to the same 150Sm abundance value on the existing abundance curve gives the means to derive the “missing” elements. This modeling with the “main” s-process yields works well for the heavier elements such as Xe. For the lighter elements such as Kr, additional contributions from the weak s-process operating in more massive stars must be considered (e.g., see reviews by Busso, Gallino, & Wasserburg, 1999; Sneden et al., 2008). The recommended abundances of Kr and Xe in Tables 4, 6, and 8 are based on interpolations with the elemental nuclide abundances resulting from the adopted compositions and s-process systematics, and from interpolation of their odd-numbered nuclides in the overall distribution curve for odd numbered nuclides.

Solar System Abundances

The proto-solar, or solar system, abundances were traditionally derived from photospheric, meteoritic, and, for some elements, theoretical considerations. Table 6 lists the solar system abundances published over time on a scale relative to 106 silicon atoms. Sometimes these abundances are referred to as “solar,” “cosmic,” or “local galactic.” In many cases the data are derived from several sources as indicated in the column headers “ph” for photosphere and “mets” for meteorite data. An “o” means other considerations than meteoritic or solar abundances were included to larger extent.

An important difference between “present” and “proto-solar” results from element diffusion and gravitational settling from the solar convection zone, which lowered the abundances of elements heavier than H (see previous discussion and Table 7). Radioactive decay of “long-lived” isotopes (which have been present since the time the oldest solids formed in the solar nebula 4.567 Ga ago) decreased abundances of the radioactive parent isotopes and increased abundances of the daughter isotopes, which has had (usually small) effects on the total elemental abundances. Table 6 therefore distinguishes the point in time the abundances are relevant (“present” or “proto-solar”). However, the element settling from the outer convection zone was only been considered since the late 1980s for He and since the late 1990s for heavy elements, whereas older compilations do not distinguish solar photospheric from proto-solar composition (except for Li and isotopes of H and He).

Table 6. Solar System Abundances Relative to Si= 106 Atoms Part 1.

Proto-Solar

Present

Present

Proto-Solar

Present

Present

Present

Present

Present

Proto-Solar

Present

Proto-Solar

Sun+CI+o

Sun+CI+o

CI-chond.

Sun+CI+o

Sun+CI+o

Sun+CI+o

CI-chond.

CI-chond.

Sun+CI+o

Sun+CI+o

Sun.

ph+CI+o

ph+CI+o

ph+CI+o

ph+CI+o

Chond.

Recomm.

±σ

Recomm.

±σ

[PLJ14]

[LPG09]

[L03]

[L03]

[PJ03]

[PB93]

[AG89]

[P88]

[M85b]

[AE82]

[Cam82]

[PSZ81]

[H79]

[SZ73]

1

H

2.52E+10

3.09E+10

[5.13E+6]

2.93E+10

2.431E+10

2.884E+10

[6.27E+6]

[6.27E+6]

2.79E+10

2.82E+10

2.71E+10

2.72E+10

2.66E+10

2.50E+10

3.10E+10

2

He

2.51E+9

1.2E+8

2.59E+9

1.2E+8

[0.601]

2.47E+9

2.343E+9

2.288E+9

2.72E+9

2.82E+9

2.60E+9

2.18E+9

1.80E+9

2.00E+9

3.10E+9

3

Li

56.9

3.4

0.339 (sun)

0.088

54.8

55.6

55.47

55.47

56.5

56.5

57.1

55.0

60.0

60.0

55

60.0

4

Be

0.637

0.044

0.637

0.044

0.638

0.612

0.7374

0.7374

0.727

0.727

0.73

0.73

0.78

1.2

0.73

0.81

5

B

18.0

1.3

18.0

1.3

18.8

18.8

17.32

17.32

16.8

21.2

21.2

6.6

24

9.0

6.6

44.0

6

C

9.12E+6

1.80E+5

9.12E+6

1.80E+5

[7.60E+5]

7.19E+6

7.079E+6

7.079E+6

[7.05E+5]

[7.05E+5]

1.01E+7

1.38E+7

1.26E+7

1.21E+7

1.11E+7

7.9E+6

1.50E+7

7

N

2.19E+6

9700

2.19E+6

9700

[5.53E+4]

2.12E+6

1.950E+6

1.950E+6

[5.97E+4]

[5.97E+4]

3.13E+6

2.75E+6

2.25E+6

2.48E+6

2.31E+6

2.1E+6

3.10E+6

8

O

1.66E+7

1.30E+5

1.66E+7

1.30E+5

[7.53E6]

1.57E+7

1.413E+7

1.413E+7

[7.64E+6]

[7.64E+6]

2.38E+7

2.29E+7

2.25E+7

2.01E+7

1.84E+7

1.7E+7

2.60E+7

9

F

1270

275

1270

275

804

804

841.1

841.1

806

[59700]

843

710

930

843

780

710

1000

10

Ne

4.37E+6

1.13E+6

4.37E+6

1.13E+6

[0.0023]

3.29E+6

2.148E+6

2.148E+6

3.44E+6

2.82E+6

3.25E+6

3.76E+6

2.60E+6

2.50E+6

1.60E+6

11

Na

57800

4700

57800

4700

5.67E+4

5.77E+4

5.751E+4

5.751E+4

5.70E+4

5.70E+4

5.74E+4

5.70E+4

5.50E+4

5.70E+4

6.00E+4

5.7E+4

6.00E+4

12

Mg

1.03E+6

4.40E+4

1.03E+6

4.40E+4

1.03E+6

1.03E+6

1.020E+6

1.020E+6

1.04E+6

1.04E+6

1.07E+6

1.01E+6

1.05E+6

1.08E+6

1.06E+6

1.01E+6

1.06E+6

13

Al

81820

6110

81820

6110

81700

8.46E+4

8.410E+4

8.410E+4

8.27E+4

8.78E+4

8.49E+4

8.00E+4

8.40E+4

8.49E+4

8.50E+4

8.00E+4

8.50E+4

14

Si

1.00E+6

3.40E+4

1.00E+6

3.40E+4

1.00E+6

1.00E+6

1.000E+6

1.000E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

15

P

8260

530

8260

530

8340

8300

8373

8373

7860

9380

10400

8580

9400

10400

6500

8580

7000

16

S

4.37E+5

2.60E+4

4.37E+5

2.60E+4

4.38E+5

4.21E+5

4.449E+5

4.449E+5

4.44E+5

1.49E+6

5.15E+5

4.80E+5

4.30E+5

5.15E+5

5.00E+5

4.8E+5

5.02E+5

17

Cl

5290

810

5290

810

5170

5170

5237

5237

5180

5180

5240

5000

4700

5240

4740

5000

5700

18

Ar

97700

25300

97700

25300

[0.0096]

92700

1.025E+5

1.025E+5

1.01E+5

1.07E+5

1.07E+5

2.00E+5

1.06E+5

1.00E+5

2.10E+5

19

K

3611

190

3606

190

3670

3660

3697

3692

3660

3660

3770

3480

3400

3770

3500

3480

3500

20

Ca

57234

4500

57239

4500

59700

60400

62870

62870

61200

431000

61100

59100

62000

61100

62500

59100

72000

21

Sc

33.7

2.2

33.7

2.2

33.9

34.4

34.20

34.20

34.5

34.5

35.5

35

35.0

33.8

31

35

31

22

Ti

2459

150

2459

150

2430

2470

2422

2422

2520

2420

2400

2420

2700

2400

2400

2420

2400

23

V

275

18

275

18

281

386

288.4

288.4

280

280

293

290

260

295

254

290

254

24

Cr

13130

460

13130

460

13200

13100

12860

12860

13400

13400

13500

13500

12900

13400

12700

13500

12700

25

Mn

9090

620

9090

620

9150

9220

9168

9168

9250

9250

9550

8720

7700

9510

9300

8720

9300

26

Fe

8.72E+5

3.80E+4

8.72E+5

3.80E+4

8.77E+5

8.48E+5

8.380E+5

8.380E+5

8.68E+5

8.58E+5

9.00E+5

8.60E+5

8.80E+5

9.00E+5

9.00E+5

8.60E+5

9.01E+5

9.00E+5

27

Co

2260

100

2260

100

2280

2350

2323

2323

2260

2260

2250

2240

2100

2250

2220

2240

2200

2300

28

Ni

48670

2940

48670

2940

4.83E+4

4.90E+4

4.780E+4

4.780E+4

4.82E+4

4.82E+4

4.93E+4

4.83E+4

4.80E+4

4.93E+4

4.78E+4

4.83E+4

4.78E+4

4.90E+4

29

Cu

535

50

535

50

549

541

527

527

542

542

522

450

520

514

540

450

540

540

30

Zn

1260

180

1260

180

1240

1300

1226

1226

1300

1300

1260

1400

980

1260

1260

1400

1260

1260

31

Ga

36.2

1.8

36.2

1.8

36.2

36.6

35.97

35.97

36.6

36.6

37.8

34.0

37.8

38

34.0

14

45.0

32

Ge

120

7

120

7

118

115

120.6

120.6

118

118

119

114

118

117

114

117

117

33

As

6.07

0.5

6.07

0.5

6.10

6.10

6.089

6.089

6.35

6.35

6.56

6.50

6.79

6.2

6.5

6.2

6.60

34

Se

67.6

5

67.6

5

67.53

67.5

65.79

65.79

71.3

70.9

62.1

63

62.1

67

63

67

66.0

35

Br

12.3

2.9

12.3

2.9

10.7

10.7

11.32

11.32

11.5

11.5

11.8

8.0

11.8

9.2

8

14

18.0

36

Kr

51.3

10.4

51.3

10.4

[1.64E–4]

55.8

55.15

55.15

45.0

45.3

41.3

40

60.0

37

Rb

7.17

0.47

7.04

0.46

7.13

7.10

6.694

6.572

7.14

7.14

7.09

6.4

7.09

6.10

6.4

6.0

6.4

38

Sr

23.3

1.3

23.4

1.3

23.3

23.4

23.52

23.64

21.8

21.8

23.5

26

23.8

22.9

26

27.0

27.0

39

Y

4.35

0.24

4.35

0.24

4.31

4.63

4.608

4.608

4.61

4.63

4.64

4.64

4.64

4.8

4.3

4.8

4.80

40

Zr

10.9

1

10.9

1

10.4

10.80

11.33

11.33

11.10

11.16

11.4

11.2

10.7

12

11

9.10

13.00

41

Nb

0.78

0.07

0.78

0.07

0.800

0.780

0.7554

0.7554

0.699

0.696

0.698

0.70

0.7

0.9

0.85

0.9

0.900

42

Mo

2.6

0.26

2.6

0.26

2.66

2.55

2.601

2.601

2.54

2.54

2.55

2.5

2.52

4

2.5

4.0

2.50

43

Tc

44

Ru

1.81

0.02

1.81

0.02

1.79

1.78

1.900

1.900

1.78

1.86

1.86

1.8

1.86

1.9

1.8

1.90

1.90

45

Rh

0.338

0.015

0.338

0.015

0.337

0.370

0.3708

0.3708

0.358

0.342

0.344

0.33

0.344

0.40

0.33

0.40

0.45

46

Pd

1.38

0.07

1.38

0.07

1.38

1.36

1.435

1.435

1.37

1.37

1.39

1.32

1.39

1.3

1.32

1.30

1.30

47

Ag

0.497

0.022

0.497

0.022

0.489

0.489

0.4913

0.4913

0.480

0.480

0.486

0.5

0.529

0.46

0.5

0.46

0.46

48

Cd

1.58

0.06

1.58

0.06

1.49

1.57

1.584

1.584

1.59

1.59

1.61

1.80

1.69

1.55

1.3

1.55

1.55

49

In

0.179

0.008

0.179

0.008

0.168

0.178

0.1810

0.1810

0.178

0.178

0.184

0.17

0.184

0.19

0.174

0.19

0.19

50

Sn

3.59

0.22

3.59

0.22

3.35

3.60

3.733

3.733

3.72

3.72

3.82

3.88

3.82

3.7

2.4

3.7

3.6

51

Sb

0.359

0.045

0.359

0.045

0.300

0.313

0.3292

0.3292

0.287

0.287

0.309

0.27

0.352

0.31

0.27

0.31

0.31

52

Te

4.72

0.23

4.72

0.23

4.56

4.69

4.815

4.815

4.68

4.68

4.81

4.83

4.91

6.5

4.8

6.5

6.5

53

I

1.59

0.64

1.59

0.64

1.05

1.10

0.9975

0.9975

0.90

0.90

0.90

1.16

0.90

1.27

1.16

1.16

1.20

54

Xe

5.50

1.11

5.50

1.11

[1.74E–4]

5.46

5.391

5.391

4.70

4.35

5.84

6.1

6.3

7.8

55

Cs

0.368

0.043

0.368

0.043

0.371

0.371

0.3671

0.3671

0.372

0.372

0.372

0.37

0.372

0.39

0.37

0.39

0.38

56

Ba

4.55

0.27

4.55

0.27

4.63

4.47

4.351

4.351

4.61

4.61

4.49

4.22

4.36

4.8

4.2

4.8

4.8

57

La

0.459

0.024

0.459

0.024

0.4561

0.457

0.4405

0.4405

0.464

0.464

0.446

0.46

0.448

0.37

0.46

0.37

0.380

58

Ce

1.16

0.05

1.16

0.05

1.160

1.18

1.169

1.169

1.20

1.20

1.14

1.2

1.16

1.20

1.2

1.20

1.20

59

Pr

0.175

0.012

0.175

0.012

0.1749

0.172

0.1737

0.1737

0.180

0.180

0.167

0.18

0.174

0.18

0.18

0.180

0.18

60

Nd

0.864

0.022

0.865

0.022

0.8621

0.857

0.8343

0.8355

0.864

0.864

0.828

0.87

0.870

0.70

0.87

0.790

0.80

61

Pm

62

Sm

0.273

0.012

0.271

0.012

0.2681

0.265

0.2542

0.2542

0.269

0.269

0.258

0.27

0.27

0.24

0.27

0.24

0.24

63

Eu

0.1

0.005

0.1

0.005

0.1016

0.0984

0.09513

0.09513

0.100

0.100

0.0973

0

0.100

0.094

0.0972

0.094

0.094

64

Gd

0.346

0.013

0.346

0.013

0.3453

0.360

0.3321

0.3321

0.341

0.341

0.330

0.34

0.34

0.42

0.331

0.12

0.42

65

Tb

0.0625

0.0025

0.0625

0.0025

0.06271

0.0634

0.05907

0.05907

0.0621

0.0621

0.0603

0.061

0.061

0.076

0.0589

0.076

0.07

66

Dy

0.407

0.016

0.407

0.016

0.4132

0.404

0.3862

0.3862

0.411

0.411

0.394

0.41

0.41

0.37

0.398

0.37

0.37

67

Ho

0.0891

0.0035

0.0891

0.0035

0.08982

0.0910

0.08986

0.08986

0.0904

0.0904

0.0889

0.091

0.091

0.092

0.0875

0.092

0.093

68

Er

0.256

0.009

0.256

0.009

0.2597

0.262

0.2554

0.2554

0.261

0.261

0.251

0.26

0.26

0.23

0.253

0.23

0.24

69

Tm

0.0403

0.0014

0.0403

0.0014

0.04054

0.0406

0.03700

0.03700

0.0399

0.0399

0.0378

0.041

0.041

0.035

0.0386

0.035

0.036

70

Yb

0.252

0.009

0.252

0.009

0.2559

0.256

0.2484

0.2484

0.251

0.251

0.248

0.25

0.25

0.20

0.243

0.20

0.22

71

Lu

0.0381

0.0018

0.038

0.0018

0.03755

0.0380

0.03580

0.03572

0.0382

0.0382

0.0367

0.038

0.038

0.035

0.0369

0.035

0.04

72

Hf

0.155

0.011

0.155

0.011

0.1566

0.156

0.1698

0.1699

0.158

0.158

0.154

0.157

0.180

0.17

0.176

0.17

0.22

73

Ta

0.0215

0.001

0.0215

0.001

0.0218

0.0210

0.02099

0.02099

0.0206

0.0203

0.0207

0.02

0.021

0.02

0.0226

0.020

0.030

74

W

0.144

0.013

0.144

0.013

0.137

0.137

0.1277

0.1277

0.129

0.136

0.133

0.13

0.127

0.30

0.137

0.30

0.16

75

Re

0.0547

0.0042

0.0521

0.0042

0.0554

0.0554

0.05509

0.05254

0.0558

0.0541

0.0517

0.052

0.052

0.051

0.0507

0.051

0.052

76

Os

0.652

0.04

0.655

0.04

0.6830

0.680

0.6713

0.6738

0.699

0.672

0.675

0.68

0.72

0.69

0.717

0.69

0.85

77

Ir

0.633

0.029

0.633

0.029

0.640

0.672

0.6448

0.6448

0.657

0.628

0.661

0.65

0.66

0.72

0.65

0.72

0.73

78

Pt

1.24

0.1

1.24

0.1

1.24

1.27

1.357

1.357

1.32

1.34

1.34

0.14

1.37

1.41

1.42

1.40

1.40

79

Au

0.195

0.016

0.195

0.016

0.197

0.195

0.1955

0.1955

0.198

0.203

0.187

0.19

0.186

0.21

0.19

0.21

0.21

80

Hg

0.376

0.156

0.376

0.156

0.41

0.458

0.4128

0.4128

0.406

0.406

0.340

0.52

0.52

0.21

0.40

1.40

0.44

81

Tl

0.179

0.015

0.179

0.015

0.184

0.182

0.1845

0.1845

0.184

0.184

0.184

0.17

0.184

0.19

0.17

0.19

0.19

82

Pb

3.31

0.2

3.33

0.2

3.32

3.33

3.234

3.258

3.21

3.21

3.15

3.09

3.15

2.6

3.1

2.6

4.0

83

Bi

0.141

0.01

0.141

0.01

0.138

0.138

0.1388

0.1388

0.140

0.140

0.144

0.14

0.144

0.14

0.136

0.14

0.14

90

Th

0.0421

0.0021

0.0336

0.0017

0.0339

0.0351

0.04399

0.03512

0.0338

0.0338

0.0335

0.032

0.0335

0.045

0.032

0.057

0.0105

92

U

0.02389

0.0017

0.00897

0.00064

0.00893

0.00893

0.024631

0.009306

0.00860

0.0086

0.0090

0.0091

0.0090

0.027

0.0091

0.0270

ph+CI+o

ph+CI+o

ph+CM2+o

ph+CI+o

Photo+chon+o

Sun

Chond.

Sun+mets+o

Sun+mets+o

Sun+mets+o

Chond.

Sun+mets+o

Sun+mets+o

Sun+mets+o

Sun+mets+o

Mets.

Mets

Sun

[Cam73a]

[U72]

[Go69]

[Cam67]

[Cam63]

[All61]

[E61]

[Cam59]

[Ran56]

[SU56]

[U52]

[S49]

[Bro49]

[S47b]

[Gs37]

[NN34]

[NN30]

[R29]

1

H

3.18E+10

2.25E+10

4.80E+10

2.60E+10

3.20E+10

2.94E+10

2.50E+10

3.50E+10

4.00E+10

1.62E+10

3.50E+10

2.00E+10

2.30E+9

[7.5E+4]

1.6E+10

2

He

2.21E+9

2.10E+9

3.90E+9

2.10E+9

5.00E+9

4.05E+9

3.80E+9

3.50E+9

3.08E+9

2.88E+9

3.50E+9

3.16E+9

9.0E+7

5.0E+8

3

Li

49.5

49.5

16.0

45.0

38

0.6 (sun)

38.0

100

61

100

100

100

100

100

66

90

5 (sun)

4

Be

0.81

0.81

0.690

7

1.000

20

16

20

16

20

19

19

15

140

3.2

5

B

350 [sic]

135

6

6.2

6

1580

24

20

24

20

23

22.9

23

38

0.1

6

C

1.18E+7

1.37E+7

1.70E+7

1.35E+7

1.66E+7

2.70E+6

9.30E+6

8.00E+6

3.50E+6

5.50E+6

8.00E+6

8.71E+6

2.25E+7

[1.77E+4]

[6300]

1.3E+6

7

N

3.74E+6

2.44E+6

4.60E+6

2.44E+6

3.00E+6

4.90E+6

2.40E+6

1.60E+7

6.60E+6

1.15E+7

1.60E+7

1.45E+7

7.50E+6

[10]

[63]

2.0E+6

8

O

2.15E+7

2.36E+7

4.40E+7

2.36E+7

2.90E+7

1.58E+7

2.50E+7

2.20E+7

2.15E+7

1.51E+7

2.20E+7

1.91E+7

7.50E+7

[3.45E+6]

[3.33e+6]

5.0E+7

9

F

2450

3630

2500

3630

10000

1600

300

1600

300

1510

9000

1510

1500

290

10

Ne

3.44E+6

2.36E+6

4.40E+6

2.36E+6

1.70E+7

8.00E+5

9E+5 to 2.4E+7

8.60E+6

1.82E+7

3.70E+6

2.29E+7

2.50E+6

11

Na

6.00E+4

5.93E+4

3.50E+4

6.32E+4

4.18E+4

7.70E+4

4.38E+4

5.10E+4

4.38E+4

5.10E+4

4.47E+4

4.62E+4

4.37E+4

4.42E+4

4.3E+4

4.00E+4

7.9E+5

12

Mg

1.061E+6

1.05E+6

1.04E+6

1.05E+6

1.05E+6

1.78E+6

9.12E+5

8.70E+5

9.12E+5

8.70E+5

8.71E+5

8.87E+5

8.71E+5

8.72E+5

8.6E+5

8.33E+5

3.2E+6

13

Al

8.50E+4

8.50E+4

8.40E+4

8.51E+4

8.93E+4

7.40E+4

9.48E+4

8.20E+4

9.48E+4

8.20E+4

8.71E+4

8.82E+4

8.71E+4

8.79E+4

8.7E+4

8.00E+4

1.3E+5

14

Si

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

1.00E+6

15

P

9600

12700

8100

12700

9320

19000

10000

7500

10000

7500

5900

13000

5900

5900

6500

7300

16

S

5.00E+5

5.05E+5

8.00E+5

5.06E+5

6.00E+5

5.20E+5

3.75E+5

9.80E+4

3.75E+5

9.80E+4

1.15E+5

3.50E+5

1.15E+5

1.14E+5

8.73E+4

1.33E+4

2.5E+4

17

Cl

5700

2000

2100

1970

1836

300000

2770

2100

8850

2100

3890

17000

3890

3900

2000

3300

18

Ar

1.17E+5

1.00E+5

3.40E+5

2.28E+5

2.40E+5

1.00E+5

1.50E+5

5.30E+4

1.50E+5

2.50E+4

1.58E+5

1.00E+5

1.26E+5

1.00E+6

19

K

4200

3200

2100

3240

2970

3900

3290

3160

3500

3160

3160

6920

6930

6900

6900

6560

8700

316000

20

Ca

72100

66600

72000

73600

72800

83000

45000

49000

56000

49000

56000

57500

67000

57500

57100

6.5E+5

63300

2.5E+5

21

Sc

35

35

35

33

29

42

32

28

21

28

17

15

18

15

15

2

320

200

22

Ti

2775

2500

2400

2300

3140

1800

2090

1680

1800

2240

1800

4680

2600

4700

4700

4700

5700

7900

23

V

262

298

590

900

590

300

220

150

220

150

129

250

130

130

270

770

5000

24

Cr

12700

11900

12400

12400

12000

1900

6400

7800

8200

7800

8200

11200

9500

11200

11300

52400

13000

25100

25

Mn

9300

9300

6200

8800

6320

5600

7200

6850

6800

6850

6800

6600

7700

6600

6600

6600

5300

40000

26

Fe

8.30E+5

8.90E+5

2.50E+5

8.90E+5

8.42E+5

4.80E+5

8.50E+5

6.70E+5

6.00E+5

6.70E+5

8.91E+5

1.83E+6

1.15E+6

8.90E+5

6.11E+5

1.97E+6

7.9E+5

27

Co

2210

2300

1900

2300

2290

2200

1190

1800

2900

1800

2900

3470

9900

3500

3500

3000

9000

20000

28

Ni

4.80E+4

4.57E+4

4.50E+4

4.57E+4

4.44E+4

4.40E+4

2.48E+4

2.74E+4

3.90E+4

2.74E+4

3.90E+4

4.57E+4

1.34E+5

4.57E+4

4.6E+4

3.9E+4

1.4E+5

5.0E+4

29

Cu

540

920

420

919

861

932

186

212

420

212

420

230

460

230

460

500

1300

5000

30

Zn

1244

1350

630

1500

930

2880

202

180

486

180

1510

160

1510

360

180

430

4000

31

Ga

48

51

28

45.5

39

2.5

9.05

10

11.4

10

62

65

66

8

10

17

5

32

Ge

115

135

76

126

134

25

18.7

25.3

110

50.5

110

186

250

324

188

65

570

50

33

As

6.60

6.40

3.80

7.2

4.4

1.70

38

4.0

38

34

480

18

19

36

6.3

0.2

34

Se

67.2

89

27

70.1

18.8

18.8

18.8

13

67.6

13

257

25

234

15

15

140

35

Br

13.5

14.5

5.4

20.6

3.95

3.95

49

13.4

49

43

42

43

43

2.6

1.7

36

Kr

46.8

54.5

25

64.4

20

42.0

270

51.3

151.0

87

190

50

37

Rb

5.88

6.5

4.1

5.95

5

4.6; 5.8

6.50

15

6.50

15

17

7.1

24

6.8

12

6.7

2.5

38

Sr

26.9

24

25

58.4

21

61.0

20

18.9

41

46

41

58

37

100

107

100

39

Y

4.80

4.60

4.70

4.60

3.60

8.90

9.7

8.90

9.7

9.8

10

9.8

9.7

9.7

50

20

40

Zr

28

26

23

30

21.00

14.20

57

54.50

140

63

150

57

140

150

140

16

41

Nb

1.4

1.3

0.9

1.2

0.810

0.81

0.7

1.00

0.70

5.60

0.90

3.89

7.0

0.16

3.7

0.5

42

Mo

4.0

2.20

2.50

2.52

2.42

2.42

5.9

2.42

5.90

26.00

19

9.55

9.5

6.5

20

1.3

43

Tc

0

0.002

44

Ru

1.9

1.85

1.83

1.60

1.58

1.20

0.87

2.1

1.49

2.10

5.10

9.30

3.63

3.6

1.9

20

2.5

45

Rh

0.40

0.29

0.33

0.33

0.26

0.27

0.15

0.71

0.214

0.71

1.40

3.5

0.96

1.3

0.71

4.3

0.16

46

Pd

1.3

1.28

1.33

1.50

1.0

0.68

1.3

0.68

1.3

5.6

3.2

3.39

2.5

1.5

15

0.63

47

Ag

0.45

0.96

0.33

0.50

0.26

0.131

0.26

1.9

0.26

0.90

2.30

2.7

1.62

3.2

0.55

5

0.5

48

Cd

1.48

2.40

1.20

2.12

0.89

0.89

2.2

0.89

2.2

8.7

3

4.79

2.6

2.2

9.7

7.90

49

In

0.189

0.217

0.100

0.217

0.11

0.0013

0.11

0.27

0.11

0.27

1.30

1.0

0.71

0.23

0.31

0.09

0.05

50

Sn

3.6

4.2

1.7

4.22

1.33

1.33

18

1.33

18

16

62

9.12

29

50

240

0.8

51

Sb

0.316

0.380

0.200

0.381

0.150

0.227

0.79

0.246

0.79

0.72

1.7

0.72

0.73

0.79

2.4

0.32

52

Te

6.42

6.80

3.10

6.76

3.0

(0.50–4.0)

2.91

0.16

4.67

0.16

14.00

?

7.94

0.2

0.09

1.7

53

I

1.09

1.40

0.41

1.41

0.46

(0.044–0.67)

0.6

1.5

0.8

1.5

1.3

1.8

8.51

1.4

0.04

54

Xe

5.38

6.62

3.00

7.10

3.15

3.35

2.5

4.0

6.3

1.5

3.72

5

55

Cs

0.387

0.370

0.210

0.367

0.25

0.100

0.456

0.1

0.456

1.3

0.4

0.1

0.35

0.1

1.3

0.10

56

Ba

4.8

4.67

5.00

4.7

4.0

3.96

3.66

9

3.66

3.3

8.3

3.9

7.76

2.1

75

19

100

57

La

0.445

0.36

0.47

0.36

0.380

0.400

0.5

2.1

2.0

2.1

2.1

2.1

2.09

2.1

2.1

3.2

58

Ce

1.18

1.17

1.38

1.17

1.08

0.62

0.58

2.3

2.26

2.3

5.2

2.3

5.25

5.2

2.3

3.67

13

59

Pr

0.149

0.17

0.19

0.17

0.160

0.150

0.23

0.96

0.40

0.96

0.95

0.96

0.95

0.96

0.97

0.2

60

Nd

0.78

0.77

0.88

0.77

0.690

0.740

0.874

3.3

1.44

3.3

3.3

3.3

3.3

3.3

3.3

2.6

5

61

Pm

62

Sm

0.226

0.23

0.28

0.23

0.24

0.25

0.238

1.1

0.664

1.1

1.1

1.2

1.1

1.2

1.2

2.6

1.6

63

Eu

0.085

0.091

0.100

0.091

0.083

0.097

0.115

0.28

0.187

0.28

0.28

0.28

0.28

0.28

0.28

1.3

64

Gd

0.297

0.40

0.43

0.34

0.33

0.36

0.516

1.6

0.684

1.6

1.7

1.7

1.7

1.65

1.6

0.63

65

Tb

0.055

0.074

0.061

0.052

0.054

0.056

0.09

0.52

0.0956

0.52

0.37

0.52

0.52

0.52

0.52

66

Dy

0.36

0.36

0.45

0.36

0.33

0.39

0.665

2

0.556

2

2

2.0

2.0

2.0

2.0

2.0

67

Ho

0.079

0.090

0.093

0.090

0.076

0.078

0.18

0.57

0.118

0.57

0.44

0.57

0.58

0.57

0.57

68

Er

0.225

0.22

0.28

0.22

0.21

0.21

0.583

1.6

0.316

1.6

1.6

1.6

1.6

1.6

1.6

0.06

69

Tm

0.034

0.035

0.041

0.035

0.032

0.039

0.09

0.29

0.0318

0.29

0.29

0.29

0.29

0.29

0.29

70

Yb

0.216

0.21

0.22

0.21

0.18

0.19

0.393

1.5

0.220

1.5

1.8

1.5

1.5

1.5

1.5

71

Lu

0.036

0.035

0.036

0.035

0.031

0.036

0.0358

0.48

0.0500

0.48

0.49

0.48

0.49

0.490

0.48

72

Hf

0.21

0.47

0.31

0.16

0.16

0.113

1.4

0.438

1.4

2.4

0.7

3.0

1.5

1.7

0.7

0.1

73

Ta

0.021

0.025

0.019

0.022

0.021

0.019

0.015

0.26

0.0650

0.26

0.40

0.31

0.45

2.8

0.07

1

74

W

0.16

0.16

0.16

0.16

0.11

0.10

0.105

13?

0.490

0.13

10.00

17

5

15

13

0.16

0.08

75

Re

0.053

0.052

0.059

0.055

0.054

0.054

0.07

0.1350

0.07

0.22

0.41

0.22

0.002

0.002

0.004

76

Os

0.75

0.70

0.86

0.71

0.73

0.60

0.64

0.97

1.000

0.97

1.70

3.5

1.7

1.7

0.72

4.7

77

Ir

0.717

0.73

0.96

0.43

0.50

0.38

0.494

0.31

0.821

0.31

1.00

1.4

1.0

0.58

0.28

1.1

0.03

78

Pt

1.40

1.30

1.40

1.13

1.16

1.28

1.5

1.63

1.5

2.9

8.7

2.9

2.9

1.8

7.7

2.0

79

Au

0.202

0.19

0.18

0.20

0.13

0.13

0.145

0.21

0.145

0.21

0.58

0.82

0.58

0.57

0.21

0.63

80

Hg

0.40

0.472

0.600

0.75

0.27

(0.076)

0.408

< 0.006

0.284

0.8

1.91

0.3

81

Tl

0.192

0.185

0.130

0.182

0.11

0.0010–0.00038

0.31

0.11

0.108

0.110

0.680

1.1

0.17

0.11

0.02

1.3

82

Pb

4

3.8

1.3

2.90

2.2

0.05–0.28

21.7

< 2

0.47

9.1

27.00

0.58

9.2

13

70

0.8

83

Bi

0.143

0.16

0.19

0.164

0.14

0.0016

0.3

0.014

0.144

0.001

0.112

0.210

0.12

0.11

0.14

0.33

90

Th

0.058

0.053

0.04

0.034

0.069

0.026

0.027

0.630

1.2

0.59

2.8

1.1

92

U

0.0262

0.0334

0.010

0.0234

0.042

0.0072

0.0078

0.3

0.26

0.23

1.3

Note: Values in square brackets are not representative for the solar system. The following all use combinations of abundances from the photosphere “Sun,” CI-chondrites “CI,” and other methods “o” (interpolation, nuclear systematics): [Cam59],[Cam63], [Cam67] = Cameron (1959), all for proto-solar. [H79] Holweger (1979); [PSZ81] Palme et al. (1981); [AE82] Anders and Ebihara (1982); [Cam82] Cameron (1982), proto solar; [P88] Palme (1988); [AG89] Anders and Grevesse (1989); [PB93] Palme and Beer (1993); [L03] Lodders (2003); [PJ05] Palme and Jones (2003); [LPG09] Lodders et al. (2009); [PLJ14] Palme et al. (2014). References to other abundances: [R29] Russell (1929), solar photosphere. [NN30] Noddack and Noddack (1930), meteoritic based on stony meteorite averages (see also Table 2). [NN34] Noddack and Noddack (1934, 1935), meteoritic based on weighted meteorite phase averages (see also Table 2). [Gs37] Goldschmidt (1937), photosphere (H,C,N,O, noble gases), meteoritic based on weighted meteorite phase averages (see also Table 2). [S47b] Suess (1947b), photosphere, meteoritic, interpolation. [Bro49] Brown (1949), photospheric (H,C,N,O, noble gases), meteoritic based on weighted meteorite phase averages to match Earth’s density (see Table 2). [S49] Suess (1949), as in [S47b]. [U52] Urey (1952), photospheric, direct chondrite compositions and weighted average of meteorite phases as observed in meteorites by Prior (1916). [SU56] Suess and Urey (1956), as in [U52], nuclear systematics and interpolation. [Ran56] Rankama (1956), photosphere, meteoritic from various sources including [Gs37], [Bro49], [U52]. [SU56], meteorites, heavier on nuclear systematics, interpolation. [All61] Aller (1961), photosphere; [E61] Ehman (1961), chondrites; [Go69] Goles (1969), photosphere, CM2-chondrites; [U72] Urey (1972), photosphere, CI- and CM2-chondrites, interpolation. [SZ73] Suess and Zeh (1973), chondrites and interpolation. [M85b] Meyer (1985a), “local galactic,” photosphere.

The calculation of the mass fractions X (for H), Y (He), and Z (for heavy elements Li-U) from the abundances on the atomic scale are done by multiplying the abundance of an element by its atomic weight and renormalizing the mass abundances to unity (or 100%). The atomic weights for several elements are different in the Sun than on Earth because of their different isotopic compositions (e.g., for C,N,O, Ar; see Table 9). Therefore, atomic weights derived from the adopted isotopic compositions of the elements were used. Calculating mass fractions requires knowledge of abundances of all (major) elements and can be done for the present-day compositions. If the He abundance is unknown, only the mass fraction ratio (Z/X) of hydrogen (X) and that of all other elements (Z) can be calculated; this ratio is an important parameter in solar models from which the He abundances are actually derived. The recommended atomic abundance of He in Table 6 and 8 comes from the helioseismically determined mass fraction of Y (Table 7).

The mass fractions are related such as X=(1Y)/[ 1+Z/X ] and Z=1XY. The atomic ratio n(He)/n(H) and the He abundance A(He)=log10{ n(He)/n(H) }+12 on the atomic scale depend on Y and Z/X because

n(He)n(H)=YX=Y(1+X/Z)4(1Y)

Where the “4” is shorthand for the ratio of the atomic weights of He and H.

The proto-solar mass fractions Xo, Yo, and Zo (indicated by subscript “0”) are derived from the present-day mass fraction by applying the settling factors SF(He) and SF(Li-U) to the present-day mass ratios Y/X and Z/X. The settling factors should never be applied to the individual fractions Y or Z. From the definition of the settling factor we have

SF(He)=(Y/X)/(Yo/Xo)

The proto-solar ratio Yo/Xo=(Y/X)/SF(He) follows from the present-day mass fraction ratio. The same applies to Zo/Xo using the settling factor SF(Li-U). With the mass-balance relations, Xo=1/(1+Yo/Xo+Zo/Xo), all proto-solar mass fractions are derived.

Note the atomic weights cancel out in the definition of the settling factors and the individual protosolar abundances on the atomic scale relative to hydrogen are easily obtained. The settling correction increases the present-day abundances on the abundance scale normalized to hydrogen by a factor “SF,” which means adding a constant log10SF in the logarithmic scale with A(H)=12:

A(E)proto-solar=12+log10{ n(E)/n(H) }proto-solar=12+A(E)present+log10SF=12+log10{ n(E)/n(H) }present+log10SF

The settling factors relate to the mass fraction ratios as log10SF(He)=log10(Y/X)/(Yo/Xo) for the helium/hydrogen ratio and log10SF(Li-U)=log10(Z/X)/(Zo/Xo) for the mass fraction ratio of all other elements. Table 7 summarizes the mass fractions X, Y, Z, their ratios, and settling factors of the recommended compositions (in Table 8) and those of other compilations and standard solar models using these compositions. The settling corrections affect the He and heavy element abundances on the logarithmic atomic scale normalized to hydrogen. However, on the cosmochemical abundance scale, only the H and He abundances are affected because this scale is relative to 106Si.

Table 7. Mass Fractions of Hydrogen (X), Helium (Y), and All Other Heavier Elements (Z) in the Solar System

X

Y

Z

Z/X

Y/X

Xo

Yo

Zo

Zo/Xo

Yo/Xo

log SF(He)

log SF(Li-U)

solar system based on:

Photosphere (3D) and CI-chondrites

≡0.7389

0.2463

0.0148

0.0200

0.3334

0.7061

0.2766

0.0173

0.0245

0.3917

≡0.07

≡0.088

Photosphere (3D)

≡0.7389

0.2462

0.0149

0.0201

0.3333

0.7061

0.2765

0.0174

0.0246

0.3916

≡0.07

≡0.088

Photosphere (1D)

≡0.7389

0.2458

0.0153

0.0207

0.3327

0.7061

0.2760

0.0179

0.0253

0.3908

≡0.07

≡0.088

Helioseismic Measurements

Convection Zone Sun (Photosphere)

Basu & Antia (2004),

0.7389±0.0035

0.2485±0.0035

0.0126

≡0.0171

0.3318

Basu & Antia (2004)

0.7389±0.0035

0.2450±0.0035

0.0161

≡0.0218

0.3316

Bahcall et al. (2006)

0.249±0.003

Asplund et al. (2009) (AGSS09) composition

0.7381

≡0.2485

0.0134

0.0181

0.3367

0.7154

0.2703

0.0142

0.0199

0.3778

≡0.05

≡0.04

Standard Solar Models based on Z/X by AGSS09

Zhang et al. (2019)

0.7484

0.2381

0.0135

≡0.0181

0.3182

0.7190

0.2662

0.0148

0.0206

0.3703

0.066

0.055

Yang (2016)

0.7507

0.2351

0.0142

≡0.0189

0.3131

0.7177

0.2664

0.0159

0.0221

0.3712

0.075

0.069

Basu et al. (2015)

0.7513

0.2352

0.0135

≡0.0180

0.3131

...

0.2650

...

...

...

...

...

Serenelli et al. (2011)

0.7547

0.2319

0.0134

≡0.0178

0.3073

0.7231

0.2620

0.0149

0.0206

0.3623

0.072

0.064

Asplund et al. (2005) (A05) composition

0..7392

≡0.2486

0.0122

0.0165

0.3363

0.7135

0.2735

0.0132

0.0185

0.3833

≡0.05

≡0.04

Standard Solar Models based on Z/X by A05

Basu et al. (2015)

0.7589

0.2286

0.0125

≡0.0165

0.3012

0.2586

Castro et al. (2007)

0.7644

0.2330

0.0125

≡0.0164

0.2917

0.7304

0.2562

0.0134

0.0183

0.3508

0.080

0.048

Bahcall et al. (2006)

0.7583

0.2291

0.0126

≡0.0166

0.3021

0.7259

0.2600

0.0141

0.0194

0.3582

0.074

0.068

Turck-Cheize et al. (2004)

0.7518

0.2353

0.0129

≡0.0172

0.3130

0.7195

0.2664

0.0141

0.0196

0.3703

0.073

0.057

Grevesse & Sauval (1998) (GS98) composition

0.7347

≡0.2483

0.0169

0.0231

0.3380

0.7086

≡0.275

0.0163

0.0230

0.3881

≡0.060

≡0

Standard Solar Models based on Z/X by GS98

Zhang et al. (2019)

0.7378

0.2453

0.0169

≡0.0229

0.03325

0.7089

0.2727

0.01837

0.0259

0.3847

0.063

0.054

Yang (2019)

0.7364

0.2461

0.0175

≡0.0238

0.3342

0.7038

0.2767

0.195

0.0277

0.3932

0.071

0.066

Basu et al. (2015)

0.7374

0.2456

0.0170

≡0.0230

0.3338

...

0.2755

...

...

...

...

...

Serenelli et al. (2011)

0.7401

0.2429

0.0170

≡0.0229

0.3282

0.7089

0.2724

0.0187

0.0264

0.3843

0.0685

0.061

Bahcall et al. (2006)

0.7404

0.2426

0.0170

≡0.02292

0.3276

0.7087

0.2725

0.01884

0.0266

0.3845

0.070

0.064

Note: Entries with ≡ are adopted quantities.

Using solar evolution models from results of Boothroyd and Sackmann (2003) and the dependence of the helium mass fraction (Y) to hydrogen (X) ratio (Z/X) and of all other elements (Z/X), Lodders (2003) derived settling factors for helium log10SF(He)=0.084 (a 22% increase from the present-day photospheric He/H ratio) and for all other heavy elements log10SFLi-U=0.074(18%). Lodders et al. (2009) and Palme et al. (2014) used 0.061 (15%) and 0.053 (13%), respectively. Asplund et al. (2005) and Asplund et al. (2009) applied log10SF(He)= 0.05(12%) and 0.04 (10%) for heavy elements from diffusion models by Turcotte and Wimmer-Schweingruber (2002); note that these are generally lower than used in several standard models (e.g., Table 8). Settling factors from standard solar models (independent of photospheric reference abundances used) for He are log10SF(He)= 0.070–0.075 (18%–19%) and for the heavy elements log10SF(Li-U)= 0.064 (16%–17%) by Bahcall, Serenelli, and Basu (2006) and Yang (2016, 2018). The standard solar model in Castro et al. (2007) log10SF(He)=0.080(20%) and 0.048(12%) for the heavy elements.

The settling corrections applied here are log10SF(He)= 0.07(17%) and log10SF(Li-U)= 0.088 (23%) for the heavier elements. Note the lower settling factor for He than for heavier elements, opposite to what was previously assumed. The settling corrections are based on models by Yang (2019) who shows that the Sun’s properties inferred from helioseismology observations (the sound-speeds, the density profiles, and the depth of the solar convection zone for the present Sun) are reproduced with increased settling efficiencies by 50% (compared to previous models) and rotational mixing for solar models with lower metallicities.

Increased settling of He from the solar convection zone below its base is almost completely counteracted by rotational forces. The settling of the heavier elements is counteracted less strongly by rotation, which is why the settling factors are now reversed from earlier estimates. The net effect of heavy element settling is an increase of the opacity in the interior below the convection zone. Yang’s (2019) models agree well with observed sound-speeds and density profiles, even better than solar models using the older Grevesse and Sauval (1998) solar composition which had been a long-time favorite before more detailed rotational mixing was introduced into solar models. The best-fit models by Yang (2019) are models with initial mass fraction of heavy elements to hydrogen ratios as obtained earlier by Lodders et al. (2009) and Caffau et al. (2011; supplemented with data from Lodders, 2009; see Table 6). More comprehensive models for element diffusion and settling from the solar convection zone and considering rotation will improve our understanding of element fractions that may become discernable with advanced spectroscopic methods and better meteoritic analyses.

Recommended Solar System Abundances

The recommended elemental solar system abundances are summarized in Table 8. They are given for the present-day and for the time of solar system formation 4.567 Ga ago (“proto-solar”) on the cosmochemical and the astronomical abundance scales. For comparison, the values for CI-chondrites and the Sun (mainly photospheric values) are listed. Since the abundances in the solar convection zone and thus in the overlying photosphere have changed over the Sun’s lifetime, the present-day observed photospheric abundances for helium and all other heavier elements needed to be readjusted with the settling factors to bring them back to proto-solar values. In addition, the abundances for elements with long-lived radioactive isotope were back-calculated to 4.567 Ga ago.

Figure 4 shows the elemental abundances as a function of atomic number. Note the logarithmic scale which spans 13 orders of magnitude. The prominence of H and He is obvious, as is the steep drop to Li, Be, and B, which are made of very fragile nuclei. The meteoritic and solar abundances of Li, Be, and B are exceptionally low when compared with their neighbors of lower and higher atomic numbers. Like the nuclides of lithium, those of beryllium and boron have low binding energies of their nuclei and are fragile in stellar interiors, and their major nucleosynthetic formation processes are relatively inefficient. Further, the Li abundance in the Sun is 170 times less than in meteorites because Li is destroyed in the Sun. The abundances of C, N, and O are high; then, with increasing atomic numbers, elemental abundances drop with a minimum at Sc, then steeply rise again to a maximum at Fe. After the “Fe-peak,” abundances drop steadily with smaller peaks at certain atomic numbers (see, e.g., Lodders & Fegley [2011] for detailed descriptions and references listed in Table 6).

Figure 4. Solar system abundances of the elements as a function of atomic number.

Table 8. Recommended Atomic Solar System Abundances: Present and Proto-Solar 4.567 Ga Ago

CI-Chondrites

Sun Outer Convection Zone

Solar System *

Solar System*

Solar System *

Solar System*

(Mainly Photosphere)

Present

Proto-Solar ,

Present

Proto-Solar ,

.

.

Si = 106

±σ

Si = 106

±σ

Si = 106

±σ

±σ

±σ

.

log(X/H)+12±σ

log(X/H)+12 ±σ

1

H

4.83E+6

4.5E+5

3.09E+10

3.09E+10

2.52E+10

s

12

12

2

He

0.599

2.59E+9

1.2E+8

2.59E+9

1.2E+8

2.51E+9

1.2E+8

s

10.924

0.02

10.994

0.02

3

Li

56.9

3.4

0.339

0.088

(0.339, sun)

0.088

56.9

3.4

3.27

0.03

3.35

0.03

4

Be

0.637

0.044

0.741

0.171

0.637

0.044

0.637

0.044

1.31

0.04

1.40

0.04

5

B

18

1.3

15.5

9.1

18

1.3

18

1.3

2.77

0.03

2.85

0.03

6

C

9.44E+5

1.79E+5

9.12E+6

1.35E+6

9.12E+6

1.8E+5

9.12E+6

1.8E+5

s

8.47

0.06

8.56

0.06

7

N

44830

9700

2.19E+6

7.0E+5

2.19E+6

9700

2.19E+6

9700

s

7.85

0.12

7.94

0.12

8

O

7.38E+6

1.3E+5

1.66E+7

2.9E+6

1.66E+7

1.3E+5

1.66E+7

1.3E+5

s

8.73

0.07

8.82

0.07

9

F

1270

275

780

600

1270

275

1270

275

4.61

0.09

4.79

0.09

10

Ne

0.00234

4.37E+6

1.13E+6

4.37E+6

1.13E+6

4.37E+6

1.13E+6

s,o

8.15

0.10

8.24

0.10

11

Na

57800

4700

50120

4840

57800

4700

57800

4700

6.27

0.03

6.36

0.03

12

Mg

1.03E+6

4.4E+4

1.12E+6

1.4E+5

1.03E+6

4.4E+4

1.03E+6

4.4E+4

7.52

0.02

7.61

0.02

13

Al

81820

6110

83180

8030

81820

6110

81820

6110

6.42

0.03

6.51

0.03

14

Si

1.00E+6

3.4E+4

1.00E+6

7E+4

1.00E+6

3.4E+4

1.00E+6

3.4E+4

7.51

0.01

7.60

0.01

15

P

8260

530

8430

810

8260

530

8260

530

5.43

0.03

5.52

0.03

16

S

4.37E+5

2.6E+4

4.37E+5

5.3E+4

4.37E+5

2.6E+4

4.37E+5

2.6E+4

7.15

0.03

7.24

0.03

17

Cl

5290

810

5495

1750

5290

810

5290

810

5.23

0.06

5.32

0.06

18

Ar

0.00959

97700

25300

97700

25300

s,o

6.50

0.10

6.59

0.10

19

K

3610

190

3685

850

3606

190

3611

190

5.07

0.02

5.16

0.02

20

Ca

57239

4500

64565

4620

57239

4500

57234

4500

6.27

0.03

6.36

0.03

21

Sc

33.7

2.2

44.7

4.3

33.7

2.2

33.7

2.2

3.04

0.03

3.13

0.03

22

Ti

2460

150

2630

254

2459

150

2459

150

4.90

0.03

4.99

0.03

23

V

275

18

240

49

275

18

275

18

3.95

0.03

4.04

0.03

24

Cr

13130

460

12880

1240

13130

460

13130

460

5.63

0.02

5.72

0.02

25

Mn

9090

620

10230

730

9090

620

9090

620

5.47

0.03

5.56

0.03

26

Fe

8.72E+5

3.8E+4

9.33E+5

9.0E+4

8.72E+5

3.8E+4

8.72E+5

3.8E+4

7.45

0.02

7.54

0.02

27

Co

2260

100

2630

320

2260

100

2260

100

4.86

0.02

4.95

0.02

28

Ni

48670

2940

48980

4730

48670

2940

48670

2940

6.20

0.03

6.29

0.03

29

Cu

535

50

468

57

535

50

535

50

4.24

0.04

4.33

0.04

30

Zn

1260

180

1120

140

1260

180

1260

180

4.61

0.60

4.70

0.60

31

Ga

36.2

1.8

32.4

3.9

36.2

1.8

36.2

1.8

3.07

0.02

3.16

0.02

32

Ge

120

7

132

23

120

7

120

7

3.59

0.03

3.68

0.03

33

As

6.07

0.50

6.07

0.50

6.07

0.50

2.29

0.03

2.38

0.03

34

Se

67.6

5.0

67.6

5.0

67.6

5.0

3.34

0.03

3.43

0.03

35

Br

12.3

2.9

12.3

2.9

12.3

2.9

2.60

0.09

2.69

0.09

36

Kr

1.63E-04

51.3

10.4

51.3

10.4

o

3.22

0.08

3.31

0.08

37

Rb

7.04

0.46

9.12

1.60

7.04

0.46

7.17

0.47

2.36

0.03

2.45

0.03

38

Sr

23.4

1.3

20.9

3.1

23.4

1.3

23.3

1.3

2.88

0.02

2.97

0.02

39

Y

4.35

0.24

5.01

0.61

4.35

0.24

4.35

0.24

2.15

0.02

2.24

0.02

40

Zr

10.9

1.0

12.5

1.84

10.9

1.0

10.9

1.0

2.55

0.04

2.64

0.04

41

Nb

0.780

0.070

0.912

0.135

0.78

0.070

0.78

0.070

1.4

0.04

1.49

0.04

42

Mo

2.6

0.26

2.34

0.54

2.6

0.26

2.6

0.26

1.92

0.04

2.01

0.04

43

Tc

44

Ru

1.81

0.02

1.74

0.35

1.81

0.02

1.81

0.02

1.77

0.01

1.86

0.01

45

Rh

0.338

0.015

0.240

0.049

0.338

0.015

0.338

0.015

1.04

0.02

1.13

0.02

46

Pd

1.38

0.07

1.10

0.16

1.38

0.07

1.38

0.07

1.65

0.02

1.74

0.02

47

Ag

0.497

0.022

0.282

0.073

0.497

0.022

0.497

0.022

1.21

0.02

1.29

0.02

48

Cd

1.58

0.06

1.82

0.75

1.58

0.06

1.58

0.06

1.71

0.02

1.80

0.02

49

In

0.179

0.008

0.195

0.114

0.179

0.008

0.179

0.008

0.76

0.02

0.85

0.02

50

Sn

3.59

0.22

3.24

0.84

3.59

0.22

3.59

0.22

2.07

0.03

2.15

0.03

51

Sb

0.359

0.045

0.359

0.045

0.359

0.045

1.06

0.05

1.15

0.05

52

Te

4.72

0.23

4.72

0.23

4.72

0.23

2.18

0.02

2.27

0.02

53

I

1.59

0.64

1.59

0.64

1.59

0.64

1.71

0.15

1.80

0.15

54

Xe

3.47E-04

5.50

1.11

5.50

1.11

o

2.25

0.08

2.34

0.08

55

Cs

0.368

0.043

0.368

0.043

0.368

0.043

1.08

0.02

1.16

0.02

56

Ba

4.55

0.27

5.5

0.92

4.55

0.27

4.55

0.27

2.17

0.02

2.26

0.02

57

La

0.459

0.024

0.398

0.038

0.459

0.024

0.459

0.024

1.17

0.02

1.26

0.02

58

Ce

1.16

0.05

1.17

0.11

1.16

0.05

1.16

0.05

1.58

0.02

1.66

0.02

59

Pr

0.175

0.012

0.162

0.016

0.175

0.012

0.175

0.012

0.75

0.01

0.84

0.01

60

Nd

0.865

0.022

0.813

0.078

0.865

0.022

0.864

0.022

1.45

0.01

1.53

0.01

61

Pm

62

Sm

0.271

0.012

0.275

0.027

0.271

0.012

0.273

0.012

0.94

0.02

1.03

0.02

63

Eu

0.100

0.005

0.102

0.010

0.1

0.005

0.1

0.005

0.51

0.02

0.60

0.02

64

Gd

0.346

0.013

0.372

0.036

0.346

0.013

0.346

0.013

1.05

0.02

1.14

0.02

65

Tb

0.0625

0.0025

0.0631

0.0163

0.0625

0.0025

0.0625

0.0025

0.31

0.02

0.39

0.02

66

Dy

0.407

0.016

0.389

0.038

0.407

0.016

0.407

0.016

1.12

0.02

1.21

0.02

67

Ho

0.0891

0.0035

0.0933

0.0269

0.0891

0.0035

0.0891

0.0035

0.46

0.02

0.55

0.02

68

Er

0.256

0.009

0.263

0.032

0.256

0.009

0.256

0.009

0.92

0.02

1.01

0.02

69

Tm

0.0403

0.0014

0.0398

0.0038

0.0403

0.0014

0.0403

0.0014

0.11

0.02

0.20

0.02

70

Yb

0.252

0.009

0.219

0.063

0.252

0.009

0.252

0.009

0.91

0.02

1.00

0.02

71

Lu

0.038

0.0018

0.0389

0.0090

0.038

0.0018

0.0381

0.0018

0.09

0.02

0.18

0.02

72

Hf

0.155

0.011

0.219

0.027

0.155

0.011

0.155

0.011

0.70

0.03

0.79

0.03

73

Ta

0.0215

0.0010

0.0215

0.0010

0.0215

0.0010

−0.16

0.02

−0.07

0.02

74

W

0.144

0.013

0.209

0.060

0.144

0.013

0.144

0.013

0.67

0.04

0.76

0.04

75

Re

0.0521

0.0042

0.0521

0.0042

0.0547

0.0042

0.23

0.03

0.34

0.03

76

Os

0.655

0.040

0.776

0.095

0.655

0.040

0.652

0.040

1.33

0.03

1.41

0.03

77

Ir

0.633

0.029

0.813

0.142

0.633

0.029

0.633

0.029

1.31

0.02

1.40

0.02

78

Pt

1.24

0.10

1.24

0.10

1.24

0.10

1.60

0.03

1.69

0.03

79

Au

0.195

0.016

0.251

0.051

0.195

0.016

0.195

0.016

0.80

0.03

0.89

0.03

80

Hg

0.376

0.156

0.376

0.156

0.376

0.156

1.08

0.15

1.17

0.15

81

Tl

0.179

0.015

0.275

0.161

0.179

0.015

0.179

0.015

0.76

0.04

0.85

0.04

82

Pb

3.33

0.20

2.57

0.52

3.33

0.20

3.31

0.20

2.03

0.03

2.12

0.03

83

Bi

0.141

0.010

0.141

0.010

0.141

0.010

0.66

0.03

0.75

0.03

90

Th

0.0336

0.0017

«/ul> 0.033

0.0336

0.0017

0.0421

0.0021

0.04

0.02

0.22

0.02

92

U

0.00897

0.00064

0.00897

0.00064

0.02389

0.00170

−0.54

0.03

−0.02

0.03

Notes:

* based on CI-chondrites except as indicated: s = from sun (Table 6, recommended 3D values), o = by other means.

corrected for radioactive decay.

photospheric values corrected element settling from the outer convection zone.

Solar System Abundances of the Isotopes

The elemental abundance features in Figure 4 are better discernable in a plot of the isotopic (or nuclide) abundances as a function of mass number, which is shown in Figure 5. This type of plot has been shown in many references indicated in Table 6 and was more influential in the development of nuclear theory and nucleosynthesis (e.g., Burbidge, Burbidge, Fowler, & Hoyle, 1957; Cameron, 1957; Suess, 1947a, 1947b) than the atomic elemental abundance plot in Figure 4 (which, however, instigated early theories, e.g., Harkins, 1917). The corresponding nuclide plot reveals the control of abundances by nuclear stability, and peaks at certain mass numbers correspond to nuclei with “magic” proton and neutron numbers.

The nuclide abundances listed in Table 9 and shown in Figure 5 were calculated from the elemental proto-solar abundances in Table 8 and the isotopic composition (atom-percent) of the elements largely taken from the review by Meija et al. (2016) and some references described later. Several updates and changes to Meija et al. are included, but not all are listed here and will be described elsewhere. Compositions for elements with long-lived isotopes were computed back to an age of 4.567 Ga, and chondritic isotopic compositions were preferred, as described in Lodders (2003) and Lodders et al. (2009).

Figure 5. Abundances of the isotopes as a function of mass number.

Data from Table 9. Even-numbered nuclides are shown in blue and odd-numbered nuclides in red.

Except for H, the noble gases, C, N, and O, the isotopic composition of materials within the solar system are relatively homogeneous on a 1% level. Isotopic variations in meteorites and planetary materials are often at the per-mil level for heavier elements, which are much smaller than the uncertainties (usually more than 5%–10%) associated with the elemental abundances. In these cases variations in the relative nuclide abundance distribution (Figure 5) are much more sensitive to elemental variations than to isotopic variations of a given element.

Exceptions with large ranges in observed isotopic compositions are H and D, because deuterium burning in the pre-main sequence Sun left the Sun essentially D-free, and D enrichments in condensed materials in meteorites and comets often stem from ion-molecule reactions. The proto-solar D/H ratio (1.97[±0.35]×105; Geiss & Gloeckler, 2003) is derived from comparing the helium isotopic measurements of Jupiter’s atmosphere with that of the solar wind and mass-balance calculations accounting for 3He produced by deuterium burning. However, direct measurements of the D/H ratio from Jupiter may not be representative according to Geiss and Gloeckler. There is a large range in D/H ratios in cometary and meteoritic materials and an increase of D/H ratios in the giant planets’ atmospheres with increasing distance from the Sun, and the “correct” D/H ratio for the solar system is debatable. Related to the D/H ratio is the helium isotopic composition. The value of 3He/4He of 1.66(±0.05)×104 of Jupiter’s atmosphere is adopted as proto-solar value (Geiss & Gloeckler, 2003). The present solar 3He/4He ratio is higher than the proto-solar value because 3He is a product of D-burning. The solar wind 3He/4He from in-situ measurements shows that the various solar wind regimes are isotopically fractionated. Further very 3He-enriched SEPs from solar flares have been observed (e.g., Mason et al., 2004). This is also seen with much higher precision in Genesis regime targets (Heber et al., 2012), and the fractionation is large for He and smaller for the heavier Ne and Ar. Note that the derivation of the proto-solar D/H ratio depends on the 3He/4He ratio of the present-day Sun, hence differences in reported proto-solar D/H values exist (see Geiss & Gloeckler, 2003; Wieler, 2016; and this encyclopedia for alternate choices).

The C, N, and O isotopic compositions in solar system materials are more variable and can reach 10 to 100 per-mil levels. Measurements of the photospheric C and O isotopic compositions have large uncertainties. The recent solar wind measurements give a light isotopic composition of the solar wind (i.e., the lower-mass isotope of an element is more abundant). Relative to Earth and most meteorites, the solar wind nitrogen is poor in 15N and has a similar 14N/15N = 440 as found for Jupiter (for more discussion, see Füri & Marty, 2015; Marty et al., 2010), which is adopted for the isotope table here. The oxygen isotopic composition of the solar wind is depleted in 17O,18O by about 7% relative to Earth and most meteoritic materials (Laming et al., 2017; McKeegan et al., 2011). The C-isotopic composition is still relatively close to the terrestrial and meteoritic values. In Table 9, the noble gases (except helium) data are in Crowther and Gilmour (2013), Heber et al. (2009), Heber et al. (2012), Pepin et al. (2012), Vogel et al. (2011), and Meshik et al. (2014). Terrestrial atmospheric noble gas isotopic compositions are probably a worse proxy to the unknown isotopic noble gas solar compositions than are the solar wind measurements. However, the solar wind isotopic compositions may not be representative for the solar system as a whole because the solar wind is fractionated in elemental composition relative to the photosphere and meteorites, and corresponding mass-dependent isotopic fractionations are observed in the solar wind. However, theoretical corrections to mass-dependent fractionations of the adopted solar wind isotopic compositions were not applied here because they are not yet well understood and vary with different authors (e.g., Heber et al., 2012; Ott, 2015). On the other hand, if the solar wind is isotopically light (as seen in C, N, O, the noble gases, and probably other elements) and were representative for the solar system as a whole, the reasons for the differences between the Sun and the planets, moons, asteroids, and comets await more explanations. More needs to be explored for the elemental and isotopic abundances of the solar system, and while great progress has been achieved in the past century, more discovery awaits.

Table 9. Abundances of the Isotopes in the Solar System Scaled to Si=106 Atoms

Z

A

Isotopic Composition Atom%

N(Elemental) Si=1E+6

N(Isotopic) Si=1E+6

Mass Fraction of All Isotopes

H

1

1

99.9980

2.52E+10

7.057E-01

H

1

2

0.00197

4.97E+5

2.781E-05

100

2.52E+10

He

2

3

0.0166

4.13E+5

3.461E-05

He

2

4

99.9834

2.49E+9

2.769E-01

100

2.49E+9

Li

3

6

7.589

4.3

7.187E-10

Li

3

7

92.411

52.6

1.025E-08

100

56.9

Be

4

9

100

0.637

0.637

1.595E-10

B

5

10

19.827

3.6

1.002E-09

B

5

11

80.173

14.4

4.405E-09

100

18.0

C

6

12

98.965

9.03E+6

3.011E-03

C

6

13

1.035

9.44E+4

3.411E-05

100

9.12E+6

N

7

14

99.774

2.18E+6

8.482E-04

N

7

15

0.226

4.94E+3

2.059E-06

100

2.19E+6

O

8

16

99.777

1.66E+7

7.377E-03

O

8

17

0.035

5.80E+3

2.739E-06

O

8

18

0.188

3.13E+4

1.565E-05

100

1.66E+7

F

9

19

100

1267

1267

6.688E-07

Ne

10

20

93.1251

4.07E+6

2.261E-03

Ne

10

21

0.2236

9.76E+3

5.693E-06

Ne

10

22

6.6513

2.90E+5

1.772E-04

100

4.37E+6

Na

11

23

100

57800

57800

3.692E-05

Mg

12

24

78.992

8.10E+5

5.398E-04

Mg

12

25

10.003

1.03E+5

7.151E-05

Mg

12

26

11.005

1.13E+5

8.158E-05

100

1.03E+6

Al

13

27

100

8.18E+4

8.18E+04

6.132E-05

Si

14

28

92.230

9.223E+5

7.167E-04

Si

14

29

4.683

4.683E+4

3.768E-05

Si

14

30

3.087

3.087E+4

2.573E-05

100

1.000E+6

P

15

31

100

8260

7.109E-06

S

16

32

95.04074

415600

3.692E-04

S

16

33

0.74869

3270

2.996E-06

S

16

34

4.19599

18300

1.727E-05

S

16

36

0.01458

64

6.396E-08

100

437300

Cl

17

35

75.7647

4010

3.896E-06

Cl

17

37

24.2353

1280

1.315E-06

100

5290

Ar

18

36

84.596

82670

8.265E-05

Ar

18

38

15.830

15030

1.582E-05

Ar

18

40

0.025

23.5

2.609E-08

100

97724

K

19

39

93.132

3363

3.638E-06

K*

19

40

0.147

5

5.552E-09

K

19

41

6.721

243

2.766E-07

100

3611

Ca

20

40

96.941

55500

6.163E-05

Ca

20

42

0.647

370

4.314E-07

Ca

20

43

0.135

77

9.191E-08

Ca

20

44

2.086

1194

1.458E-06

Ca

20

46

0.004

2

2.554E-09

Ca

20

48

0.187

107

1.426E-07

100

57200

Sc

21

45

100

33.7

3.370E+01

4.209E-08

Ti

22

46

8.249

203

2.592E-07

Ti

22

47

7.437

183

2.387E-07

Ti

22

48

73.72

1810

2.411E-06

Ti

22

49

5.409

133

1.809E-07

Ti

22

50

5.185

128

1.776E-07

100

2460

V

23

50

0.2497

0.7

9.715E-10

V

23

51

99.7503

274.5

3.886E-07

100

275.2

Cr

24

50

4.3452

571

7.924E-07

Cr

24

52

83.7895

11000

1.587E-05

Cr

24

53

9.5006

1250

1.839E-06

Cr

24

54

2.3647

310

4.646E-07

100

13100

Mn

25

55

100

9090

9090

1.388E-05

Fe

26

54

5.845

51000

7.643E-05

Fe

26

56

91.754

8.01E+5

1.245E-03

Fe

26

57

2.1191

18500

2.927E-05

Fe

26

58

0.2819

2460

3.960E-06

100

8.73E+5

Co

27

59

100

2260

2260

3.701E-06

Ni

28

58

68.0769

33100

5.328E-05

Ni

28

60

26.2231

12800

2.131E-05

Ni

28

61

1.1399

555

9.396E-07

Ni

28

62

3.6345

1770

3.046E-06

Ni

28

64

0.9256

450

7.993E-07

100

48700

Cu

29

63

69.174

370

6.469E-07

Cu

29

65

30.826

165

2.977E-07

100

535

Zn

30

64

49.1704

620

1.101E-06

Zn

30

66

27.7306

349

6.393E-07

Zn

30

67

4.0401

51

9.484E-08

Zn

30

68

18.4483

233

4.397E-07

Zn

30

70

0.6106

8

1.554E-08

100

1260

Ga

31

69

60.108

21.7

4.156E-08

Ga

31

71

39.892

14.4

2.838E-08

100

36.2

Ge

32

70

20.526

24.7

4.799E-08

Ge

32

72

27.446

33.0

6.595E-08

Ge

32

73

7.76

9.3

1.884E-08

Ge

32

74

36.523

43.9

9.017E-08

Ge

32

76

7.745

9.3

1.962E-08

100

120

As

33

75

100

6.07

6.07

1.264E-08

Se

34

74

0.86

0.58

1.191E-09

Se

34

76

9.22

6.23

1.314E-08

Se

34

77

7.59

5.13

1.096E-08

Se

34

78

23.69

16.00

3.464E-08

Se

34

80

49.81

33.66

7.474E-08

Se

34

82

8.83

5.96

1.357E-08

100

67.6

Br

35

79

50.686

6.26

1.373E-08

Br

35

81

49.314

6.09

1.369E-08

100

12.3

Kr

36

78

0.362

0.19

4.114E-10

Kr

36

80

2.326

1.20

2.665E-09

Kr

36

82

11.655

5.99

1.363E-08

Kr

36

83

11.546

5.94

1.368E-08

Kr

36

84

56.903

29.18

6.803E-08

Kr

36

86

17.208

8.79

2.098E-08

100

51.3

Rb

37

85

70.844

5.080

1.199E-08

Rb*

37

87

29.156

2.091

5.049E-09

100

7.17

Sr

38

84

0.5584

0.13

3.031E-10

Sr

38

86

9.8708

2.30

5.490E-09

Sr

38

87

6.8982

1.61

3.888E-09

Sr

38

88

82.6725

19.2

4.690E-08

100

23.3

Y

39

89

100

4.35

4.35

1.075E-08

Zr

40

90

51.452

5.621

1.404E-08

Zr

40

91

11.223

1.226

3.097E-09

Zr

40

92

17.146

1.873

4.783E-09

Zr

40

94

17.38

1.899

4.955E-09

Zr^

40

96

2.799

0.306

8.154E-10

100

10.92

Nb

41

93

100

0.780

0.780

2.014E-09

Mo

42

92

14.525

0.380

9.704E-10

Mo

42

94

9.151

0.238

6.210E-10

Mo

42

95

15.838

0.412

1.086E-09

Mo

42

96

16.672

0.433

1.154E-09

Mo

42

97

9.599

0.249

6.704E-10

Mo

42

98

24.391

0.630

1.714E-09

Mo^

42

100

9.824

0.253

7.023E-10

100

2.60

Ru

44

96

5.542

0.100

2.665E-10

Ru

44

98

1.869

0.034

9.249E-11

Ru

44

99

12.758

0.230

6.321E-10

Ru

44

100

12.599

0.228

6.329E-10

Ru

44

101

17.060

0.308

8.635E-10

Ru

44

102

31.552

0.570

1.614E-09

Ru

44

104

18.621

0.336

9.700E-10

100

1.81

Rh

45

103

100

0.338

0.338

9.664E-10

Pd

46

102

1.02

0.0141

3.992E-11

Pd

46

104

11.14

0.1536

4.434E-10

Pd

46

105

22.33

0.3079

8.975E-10

Pd

46

106

27.33

0.377

1.109E-09

Pd

46

108

26.46

0.365

1.094E-09

Pd

46

110

11.72

0.162

4.947E-10

100

1.38

Ag

47

107

51.839

0.258

7.664E-10

Ag

47

109

48.161

0.239

7.232E-10

100

0.497

Cd

48

106

1.249

0.020

5.885E-11

Cd

48

108

0.89

0.014

4.197E-11

Cd

48

110

12.485

0.197

6.016E-10

Cd

48

111

12.804

0.202

6.225E-10

Cd

48

112

24.117

0.381

1.185E-09

Cd^

48

113

12.225

0.193

6.055E-10

Cd

48

114

28.729

0.454

1.437E-09

Cd^

48

116

7.501

0.119

3.832E-10

100

1.58

In

49

113

4.281

0.008

2.510E-11

In^

49

115

95.719

0.171

5.459E-10

100

0.179

Sn

50

112

0.971

0.035

1.088E-10

Sn

50

114

0.659

0.024

7.596E-11

Sn

50

115

0.339

0.012

3.831E-11

Sn

50

116

14.536

0.522

1.681E-09

Sn

50

117

7.676

0.276

8.965E-10

Sn

50

118

24.223

0.870

2.850E-09

Sn

50

119

8.585

0.308

1.018E-09

Sn

50

120

32.593

1.171

3.901E-09

Sn

50

122

4.629

0.166

5.623E-10

Sn

50

124

5.789

0.208

7.161E-10

100

3.59

Sb

51

121

57.213

0.205

6.887E-10

Sb

51

123

42.787

0.154

5.259E-10

100

0.359

Te

52

120

0.096

0.005

1.666E-11

Te

52

122

2.603

0.123

4.166E-10

Te^

52

123

0.908

0.043

1.468E-10

Te

52

124

4.816

0.227

7.815E-10

Te

52

125

7.139

0.337

1.170E-09

Te

52

126

18.952

0.894

3.127E-09

Te^

52

128

31.687

1.494

5.309E-09

Te^

52

130

33.799

1.594

5.753E-09

100

4.72

I

53

127

100

1.59

1.59

5.606E-09

Xe^

54

124

0.129

0.007

2.410E-11

Xe

54

126

0.110

0.006

2.099E-11

Xe

54

128

2.220

0.122

4.336E-10

Xe

54

129

27.428

1.507

5.398E-09

Xe

54

130

4.349

0.239

8.626E-10

Xe

54

131

21.763

1.196

4.350E-09

Xe

54

132

26.360

1.449

5.311E-09

Xe

54

134

9.730

0.535

1.991E-09

Xe^

54

136

7.911

0.435

1.643E-09

100

5.495

Cs

55

133

100

0.368

0.368

1.359E-09

Ba^

56

130

0.106

0.005

1.805E-11

Ba

56

132

0.101

0.005

1.833E-11

Ba

56

134

2.417

0.110

4.093E-10

Ba

56

135

6.592

0.300

1.125E-09

Ba

56

136

7.853

0.357

1.348E-09

Ba

56

137

11.232

0.511

1.944E-09

Ba

56

138

71.699

3.264

1.251E-08

100

4.552

La*

57

138

0.092

4.000E-04

1.533E-12

La

57

139

99.908

0.459

1.772E-09

100

0.459

Ce

58

136

0.186

0.002

7.552E-12

Ce^

58

138

0.250

0.003

1.150E-11

Ce

58

140

88.450

1.030

4.004E-09

Ce^

58

142

11.114

0.129

5.086E-10

100

1.165

Pr

59

141

100

0.175

0.175

6.851E-10

Nd

60

142

27.045

0.234

9.226E-10

Nd

60

143

12.021

0.104

4.130E-10

Nd^

60

144

23.729

0.205

8.197E-10

Nd

60

145

8.763

0.076

3.060E-10

Nd

60

146

17.130

0.148

6.000E-10

Nd

60

148

5.716

0.049

2.014E-10

Nd^

60

150

5.596

0.048

1.999E-10

100

0.864

Sm

62

144

3.083

0.008

3.199E-11

Sm*

62

147

15.017

0.042

1.714E-10

Sm^

62

148

11.254

0.031

1.274E-10

Sm

62

149

13.830

0.038

1.572E-10

Sm

62

150

7.351

0.020

8.331E-11

Sm

62

152

26.735

0.073

3.081E-10

Sm

62

154

22.730

0.062

2.652E-10

100

0.273

Eu^

63

151

47.81

0.0479

2.009E-10

Eu

63

153

52.19

0.0523

2.222E-10

100

0.1002

Gd^

64

152

0.2029

0.00013

5.357E-13

Gd

64

154

2.1809

0.00755

3.227E-11

Gd

64

155

14.7998

0.0512

2.204E-10

Gd

64

156

20.4664

0.0708

3.068E-10

Gd

64

157

15.6518

0.0542

2.361E-10

Gd

64

158

24.8347

0.0859

3.770E-10

Gd

64

160

21.8635

0.0756

3.361E-10

100

0.3454

Tb

65

159

100

0.0625

6.250E-02

2.760E-10

Dy

66

156

0.054

0.00005

2.166E-13

Dy

66

158

0.095

0.0004

1.755E-12

Dy

66

160

2.329

0.0095

4.221E-11

Dy

66

161

18.889

0.0769

3.438E-10

Dy

66

162

25.479

0.1038

4.670E-10

Dy

66

163

24.895

0.1014

4.590E-10

Dy

66

164

28.260

0.1151

5.243E-10

100

0.407

Ho

67

165

100

0.0891

8.910E-02

4.083E-10

Er

68

162

0.139

0.0004

1.800E-12

Er

68

164

1.601

0.0041

1.867E-11

Er

68

166

33.503

0.0859

3.965E-10

Er

68

167

22.869

0.0587

2.737E-10

Er

68

168

26.978

0.0692

3.220E-10

Er

68

170

14.910

0.0382

1.794E-10

100

0.256

Tm

69

169

100

0.0403

4.030E-02

1.892E-10

Yb

70

168

0.12

0.0003

1.400E-12

Yb

70

170

2.98

0.0075

3.541E-11

Yb

70

171

14.09

0.0356

1.691E-10

Yb

70

172

21.69

0.0547

2.613E-10

Yb

70

173

16.10

0.0406

1.951E-10

Yb

70

174

32.03

0.0808

3.905E-10

Yb

70

176

13.00

0.0328

1.603E-10

100

0.252

Lu

71

175

97.1795

0.0370

1.798E-10

Lu*

71

176

2.8205

0.0011

5.377E-12

100

0.0381

Hf^

72

174

0.161

0.0003

1.450E-12

Hf

72

176

5.205

0.0081

3.960E-11

Hf

72

177

18.604

0.0289

1.421E-10

Hf

72

178

27.297

0.0424

2.096E-10

Hf

72

179

13.627

0.0212

1.054E-10

Hf

72

180

35.105

0.0545

2.725E-10

100

0.155

Ta*

73

180

0.0120

2.600E-06

1.300E-14

Ta

73

181

99.9880

0.0215

1.081E-10

100

0.0215

W^

74

180

0.120

0.0002

1.000E-12

W

74

182

26.499

0.0381

1.926E-10

W

74

183

14.314

0.0206

1.047E-10

W

74

184

30.642

0.0440

2.249E-10

W

74

186

28.426

0.0408

2.108E-10

100

0.144

Re

75

185

35.662

0.0195

1.002E-10

Re*

75

187

64.338

0.0352

1.828E-10

100

0.0547

Os^

76

184

0.020

0.0001

5.111E-13

Os^

76

186

1.597

0.0104

5.373E-11

Os

76

187

1.298

0.0085

4.415E-11

Os

76

188

13.325

0.0869

4.538E-10

Os

76

189

16.252

0.106

5.565E-10

Os

76

190

26.433

0.172

9.078E-10

Os

76

192

41.076

0.268

1.429E-09

100

0.652

Ir

77

191

37.272

0.236

1.252E-09

Ir

77

193

62.728

0.397

2.129E-09

100

0.633

Pt*

78

190

0.013

0.0004

1.056E-12

Pt

78

192

0.794

0.01

5.334E-11

Pt

78

194

32.808

0.407

2.193E-09

Pt

78

195

33.787

0.419

2.270E-09

Pt

78

196

25.290

0.314

1.710E-09

Pt

78

198

7.308

0.091

5.006E-10

100

1.24

Au

79

197

100

0.195

0.195

1.067E-09

Hg

80

196

0.16

0.001

5.445E-12

Hg

80

198

10.04

0.038

2.090E-10

Hg

80

199

16.94

0.064

3.538E-10

Hg

80

200

23.14

0.087

4.834E-10

Hg

80

201

13.17

0.049

2.736E-10

Hg

80

202

29.74

0.112

6.285E-10

Hg

80

204

6.82

0.026

1.474E-10

100

0.376

Tl

81

203

29.524

0.053

2.989E-10

Tl

81

205

70.476

0.126

7.176E-10

100

0.179

Pb^

82

204

1.997

0.066

3.740E-10

Pb

82

206

18.582

0.615

3.520E-09

Pb

82

207

20.563

0.680

3.911E-09

Pb

82

208

58.858

1.947

1.125E-08

100

3.308

Bi^

83

209

100

0.1414

1.414E-01

8.210E-10

Th*

90

232

100

0.0421

4.210E-02

2.714E-10

U*

92

234

0.004

9.90E-07

6.438E-15

U*

92

235

24.302

0.0058

3.788E-11

U*

92

238

75.694

0.0181

1.197E-10

100

0.0239

Note: Entries marked with * involve long-lived radioactive nuclides with half-lives up to 1012 years and abundances are for 4.567 Ga ago. Isotopes with half-lives above 1012 years (marked with ^) can be regarded as stable compared to the age of the solar system and are of interest of studies of double-beta decay. Isotopic compositions mainly adopted from Meija et al. (2016); see text for H, He, C, N, O, Ne, Ar, Kr, and Xe isotopic compositions.

Acknowledgments

I thank Bruce Fegley, Ryan Ogliore, and Rainer Wieler for thoughtful comments and suggestions for this article. I thank Don Burnett for communications about the Genesis data. Work supported in part by NSF AST 1517541.

References

  • Alexeeva, S. A., & Mashonkina, L. I. (2015). Carbon abundances of reference late-type stars from 1D analysis of atomic C I and molecular CH lines. Monthly Notices of the Royal Astronomical Society, 453, 1619–1631.
  • Allende Prieto, C., Lambert, D.L., Asplund, M., (2001). The forbidden abundance of oxygen in the Sun. Astrophysical Journal, 556, L63-L66.
  • Allende Prieto, C. (2016). Solar and stellar photospheric abundances. Living Reviews in Solar Physics, 13, id1.
  • Aller, L. H. (1961). The abundances of the elements. New York, NY: Interscience.
  • Amarsi, A. M., & Asplund, M. (2017). The solar silicon abundance based on 3D non-LTE calculations. Monthly Notices of the Royal Astronomical Society, 464, 264–273.
  • Amarsi, A. M., Barklem, P. S., Asplund, M., Collet, R., & Zatarinny, O. (2018). Inelastic O+H collisions and the O I 777 nm solar centre-to-limb variation. Astronomy & Astrophysics, 616, A89.
  • Amarsi, A. M., Barklem, P. S., Collet, R., Grevesse, N., & Asplund, M. (2019). 3D non-LTE line formation of neutral carbon in the Sun. Astronomy & Astrophysics, 624, A111.
  • Anders, E. (1964). Origin, age, and composition of meteorites. Space Science Reviews, 3, 583–714.
  • Anders, E. (1971). How well do we know “cosmic” abundances? Geochimica et Cosmochimica Acta, 35, 516–522.
  • Anders, E., & Ebihara, M. (1982). Solar-System abundances of the elements. Geochimica et Cosmochimica Acta, 46, 2363–2380.
  • Anders, E., & Grevesse, N. (1989). Abundances of the elements: Meteoritic and solar. Geochimica et Cosmochimica Acta, 53, 197–214.
  • Arlandini, C., Käppeler, F., Wisshak, K., Gallino, R., Lugaro, M., Busso, M., & Straniero, O. (1999). Neutron capture in low mass Asymptotic Giant Branch stars: Cross sections and abundance signatures. Astrophysical Journal, 525, 886–900.
  • Asplund, M., Grevesse, N., & Sauval, A. J. (2005). The solar chemical composition. In T. G. Barnes & F. N. Bash (Eds.), Cosmic abundances as records of stellar evolution and nucleosynthesis (pp. 25–38). ASP Conference Series 336, 2538. San Francisco, CA: Astronomical Society of the Pacific.
  • Asplund, M., Grevesse, N., Sauval, A. J., & Scott, P. (2009). The chemical composition of the Sun. Annual Review of Astronomy and Astrophysics, 47, 481–522.
  • Bahcall, J. N., Serenelli, A. M., & Basu, S. (2006). 10,000 standard solar models: A Monte Carlo simulation. Astrophysical Journal, 165, 400–431.
  • Baschek, B., Garz, T., Holweger, H., & Richter, J. (1970). Experimental oscillator strengths of Fe II lines and the solar iron abundance. Astronomy & Astrophysics, 4, 229–233.
  • Basu, S., & Antia, H. M. (2004). Constraining solar abundances using helioseismology. Astrophysical Journal, 606, L85–L88.
  • Basu, S., & Antia, H.M. (2005). Helioseismology and solar abundances. Physics Reports, 457, 217-283.
  • Basu, S., & Antia, H. M. (2008). Helioseismology and solar abundances. Physics Reports, 457, 217–283.
  • Basu, S., Grevesse, N., Mathis, S., & Turck-Chieze, S. (2015). Understanding the internal chemical composition and physical processes of the solar interior. Space Science Reviews, 196, 49–77.
  • Bergemann, M., Collet, R., Amarsi, A. M., Kovalev, M., Ruchti, G., & Magic, Z. (2017). Non-local thermodynamic equilibrium stellar spectroscopy with 1D and <3D> models. I. Methods and application to magnesium abundances in standard stars. Astrophysical Journal, 847, id15.
  • Bergemann, M., Gallagher, A. J., Eitner, P., Bautista, M., Collet, R., Yakovleva, S. A., . . . Hansen, C. (in press). On the origin of the elements, I. 3D NLTE formation of Mn lines in late-type stars. Astronomy & Astrophysics.
  • Berzelius, J. (1834). Über Meteorsteine. Poggendorfs Annalen der Physik Chemie, 33, 115–120, 143.
  • Biermann, L. (1951). Kometenschweife und solare Korpuskularstrahlung. Zeitschrift fur Astrophysik, 29, 274–286.
  • Bisterzo, S., Travaglio, C., Wiescher, M., Käppeler, F., & Gallino, R. (2017). Galactic chemical evolution: The impact of the 13C pocket structure on the s-process distribution. Astrophysical Journal, 835, 97.
  • Bochsler, P. (2007). Minor ions in the solar wind. Astronomy and Astrophysical Review, 14, 1–40.
  • Bochsler, P. (2009). Composition of matter in the heliosphere, in Universal Heliophysical Processes, Proceedings of the International Astronomical Union, IAU Symposium, 257, 17–28.
  • Boothroyd, A. I., & Sackmann, I. J. (2003). Our Sun. IV. The standard model and helioseismology: Consequences of uncertainties in input physics and in observed solar parameters. Astrophysical Journal, 583, 1004–1023.
  • Brown, H. (1949). A table of relative abundances of nuclear species. Reviews in Modern Physics, 21, 625–634.
  • Buddhue, J. D. (1946). The abundances of the chemical elements in meteorites and tektites. Popular Astronomy, 54, 308–311.
  • Burbidge, E. M., Burbidge, G. R., Fowler, W. A., & Hoyle, F. (1957). Synthesis of the elements in stars. Reviews in Modern Physics, 29, 547–650.
  • Burnett, D. S., & Genesis Science Team. (2011). Solar composition from the Genesis Discovery Mission. Proceedings of the National Academy of Sciences, 108, 19147–19151.
  • Burnett, D. S., Guan, Y., Heber, V. S., Hervig, R., Huss, G. R., Jurewicz, A. J. G., . . . Woolum, D. S. (2017). Solar nebula composition based on solar wind data. Lunar Planetary Science Conference 48, #1532. Houston, TX: Lunar and Planetary Institute.
  • Burnett, D. S., Jurewicz, A. J. G., & Woolum, D. S. (2019). The future of Genesis science. Meteoritics & Planetary Science, 54, 1092–1114.
  • Busso, M., Gallino, R., & Wasserburg, G. J. (1999). Nucleosynthesis in asymptotic giant branch stars: Relevance for galactic enrichment and solar system formation. Annual Review of Astronomy and Astrophysics, 37, 239–309.
  • Caffau, E., Ludwig, H. G., Steffen, M., Freytag, B., & Bonifacio, P. (2011). Solar chemical abundances determined with a CO5BOLD 3D model atmosphere. Solar Physics, 268, 255–269.
  • Caffau, E., Ludwig, H.-G., Steffen, M., Livingston, W., Bonifacio, P., Malherbe, J.-M., . . . Schmidt, W. (2015). The photospheric solar oxygen project. III. Investigation of the centre-to-limb variation of the 630 nm [O I]-Ni I blend. Astronomy & Astrophysics, 579, A88.
  • Cameron, A. G. W. (1957). Nuclear reactions in stars and nucleogenesis. Publications of the Astronomical Society of the Pacific, 69, 201–222.
  • Cameron, A. G. W. (1959). A revised table of abundances of the elements. Astrophysical Journal, 129, 676–699.
  • Cameron, A.G.W. (1963). Nuclear Astrophysiscs. Notes from lectures at Yale Univ., quoted in Cameron (1967).
  • Cameron, A. G. W. (1968). A new table of abundances of the elements in the solar system. In L. H. Ahrens (Ed.), Origin and distribution of the elements (pp. 125–143). Oxford, UK: Pergamon Press.
  • Cameron, A. G. W. (1973). Abundances of the elements in the solar system. Space Science Reviews, 15, 121–146.
  • Cameron, A. G. W. (1982). Elemental and nuclidic abundances in the solar system. In C. A. Barnes, D. D. Clayton, & D. N. Schramm (Eds.), Essays in nuclear astrophysics (pp. 23–43). Cambridge, UK: Cambridge University Press.
  • Castro, M., Vauclair, S., & Richard, O. (2007). Low abundances of heavy elements in the solar outer layers: comparisons of solar models with helioseismic inversions. Astronomy & Astrophysics, 463, 755–758.
  • Chladni, E. F. F. (1794) Ueber den Ursprung der von Pallas gefundenen und anderer ihr ähnlicher Eisenmassen, J. F Hartknoch Publisher, Riga. (About the origin of the iron mass found by Pallas and similar ones to it, and on some connected natural phenomena, see also Journal of Mines, 15, 286–320, 466–485.)
  • Chaldni, E. F. F. (1819). Über Feuermeteore und die mit ihnen heruntergefallenen Massen. Wien, Germany: J. B. Heubner.
  • Claas, W. J. (1951). The composition of the solar atmosphere. Recherches astronomiques de l’Observatoire d’Utrecht 12, part 1. Haarlem, Holland: Enschede.
  • Clarke, F. W., & Washington, H. S. (1922). The average chemical composition of igneous rocks. Proceedings of the National Academy of Sciences, 8, 108–114.
  • Crowther, S. A., & Gilmour, J. D. (2013). The Genesis solar wind composition and its relationship to planetary xenon signatures. Geochimica et Cosmochimica Acta, 123, 17–34.
  • Dahlburg, R. B., Laming, J. M., Taylor, B., & Obenschain, K. (2016). Ponderomotive acceleration in coronal loops. Astrophysical Journal, 831, id160.
  • Daubree, G. A. (1864). Note sure les meteorites tombees le 14 mai aux environs d’Orgueil (Tarne-et-Garonne). Comptes rendus de l’Académie des Sciences, 58, 984–986, 1065–1072.
  • de Beaumont, E. (1947). Note sur les emanations volcaniques et metalliferes. Bulletin de la Société Géologique de France, IIe(Series 4), 1249.
  • Dennis, B. R., Phillips, K. J. H., Schwartz, R. A., Tolbert, A. K., Starr, R. D., & Nittler, L. R. (2015). Solar flare element abundances from the Solar Assembly for X-Rays (SAX) on MESSENGER. Astrophysical Journal, 803, id67.
  • Desai, M., & Giacalone, J. (2016). Large gradual solar energetic particle evens. Living Reviews in Solar Physics, 13, id3.
  • Eddington, A. S. (1926). The internal constitution of the stars. Cambridge, UK: Cambridge University Press.
  • Ehman, W. D. (1961). Recent improvement in our knowledge of cosmic abundances. Journal of Chemical Education, 38, 53–57.
  • Engvold, O. (1977). The solar chemical composition. Physica Scripta, 16, 48–50.
  • Farrington, O. C. (1915). Meteorites. Chicago, IL: Lake Side Press.
  • Feldman, U., & Widing, K. G. (2003). Elemental abundances in the solar upper atmosphere derived by spectroscopic means. Space Science Reviews, 107, 665–720.
  • Füri, E., & Marty, B. (2015). Nitrogen isotope variations in the solar system. Nature Geoscience, 8, 515–522.
  • Garz, T., Holweger, H., Kock, M., & Richter, J. (1969). Revised solar iron abundance and its influence on the photospheric model. Astronomy & Astrophysics, 2, 446–450.
  • Garz, T., & Kock, M. (1969). Experimental Oscillator Strengths for Fe i Lines. Astronomy &Astrophysics, 274.
  • Garz, T., Kock, M., Richter, J., Baschek, B., Holweger, H., & Unsöld, A. (1969). Abundances of iron and some other elements in the Sun and in meteorites. Nature, 223, 1254–1255.
  • Geiss, J. (1998). Constraints on the FIP mechanisms from solar wind abundance data. Space Science Reviews, 85, 241–252.
  • Geiss, J., & Bochsler, P. (1986). Ion composition in the solar wind in relation to solar abundances. In Conference proceedings: Isotopic ratios in the solar system (pp. 213–228). Paris, France: National Centre for Space Studies.
  • Geiss, J., Buehler, F., Cerutti, H., Eberhardt, P., Filleux, C., Meister, J., & Signer, P. (2004). The Apollo SWV experiment: Results, conclusions, consequences. Space Science Reviews, 110, 307–335.
  • Geiss, J., & Gloeckler, G. (2003). Isotopic composition of H, He and Ne in the protosolar cloud. Space Science Reviews, 106, 3–18.
  • Geiss, J., Gloeckler, G., & von Steiger, R. (1994). Solar and heliospheric processes from solar wind composition measurements. Philosophical Transactions of the Royal Society A, 349, 213–226.
  • Giammanco, C., Wurz, P., & Karrer, R. (2008). Minor ions abundances in the slow solar wind. Astrophysical Journal, 681, 1703–1707.
  • Gloeckler, G., Cain, J., Ipavich, F. M., Tums, E. O., Bedini, P., Fisk, L. A., . . . Kallenbach, R. (1998). Investigation of the composition of solar and interstellar matter using solar wind and pickup ion measurements with SWICS and SWIMS on the ACE spacecraft. Space Science Reviews, 86, 497–539.
  • Gloeckler, G., & Geiss, J. (2007). The composition of the solar wind in polar coronal holes. Space Science Reviews, 130, 139–152.
  • Goldberg, L., & Aller, L. H. (1943). Atoms, stars, and nebulae. In H. Shapley & B. J. Bok (Eds.), The Harvard books on astronomy. Philadelphia, PA: Blakiston.
  • Goldberg, L., Müller, E. A., & Aller, L. H. (1960). The abundances of the elements in the solar atmosphere. Astrophysical Journal Supplement Series 5. Chicago, IL: University of Chicago Press.
  • Goldschmidt, V. M. (1923). Geochemische Verteilungsgesetze der Elemente, I. Skrifter Norske Videnskaps-Akad. Akademi I. Mathematisk-Naturvidenskapelig klasse No. 3. Oslo, Norway: Norske videnskaps-akademi.
  • Goldschmidt, V. M. (1937). Geochemische Verteilungsgestze der Elemente, IX. Die Mengenverhältnisse der Elemente und der Atom-Arten, Skrifter Norske Videnskaps-Akad. Akademi I. Mathematisk-Naturvidenskapelig klasse, No. 4. Oslo, Norway: Norske videnskaps-akademi.
  • Goldschmidt, V. M. (1938). The principles of distribution of chemical elements in rocks and minerals. Journal of the Chemical Society, 140, 653–673.
  • Goles, G. G. (1969). Cosmic abundances, nucleosynthesis, and cosmic chronology. Handbook of Geochemistry, 1, 116–131.
  • Greenland, L., & Lovering, J. F. (1965). Minor and trace element abundances in chondritic meteorites. Geochimica et Cosmochimica Acta, 29, 821–833.
  • Grevesse, N. (1984). Abundances of the elements in the Sun. In R. Pallavicini (Ed.), Frontiers of astronomy and astrophysics: Invited papers presented at the 7th European Regional Meeting in Astronomy, held in Florence, Italy, December 12–16, 1983 (pp. 71–82). Florence, Italy: Italian Astronomical Society.
  • Grevesse, N., Asplund, M., & Sauval, A. J. (2007). The solar chemical composition. Space Science Reviews, 130, 105–114.
  • Grevesse, N., Blanquest, G., & Boury, A. (1968). Abundances solaires de quelques éléments représentatifs au point de vue de la nucléosynthèse, In L. H. Ahrens (Ed.), Origin and distribution of the elements (pp. 177–182). Oxford, UK: Pergamon Press.
  • Grevesse, N., & Noels, A. (1993). La composition chimique du Soleil. In N. Prantzos, E. Vangioni-Flam, & M. Cassé (Eds.), Origin and evolution of the elements (pp. 205–257). Cours de perfectionnement de l’Association Vaudoise des Chercheurs en Physique. Cambridge, UK: Cambridge University Press.
  • Grevesse, N., Noels, A., & Sauval, A. J. (1996). Standard abundances. In S. S. Holt & G. Sonneborn (Eds.), Cosmic abundances (pp. 117–126). Astronomical Society of the Pacific Conference Series 99. San Francisco, CA: Astronomical Society of the Pacific.
  • Grevesse, N., & Sauval, A. J. (1998). Standard solar composition. Space Science Reviews, 85, 161–174.
  • Grevesse, N., & Sauval, A. J. (1998). Standard solar composition. Advances in Space Research, 30, 3–11.
  • Grevesse, N., Scott, P., Asplund, M., & Sauval, A. J. (2015). The elemental composition of the Sun. III. The heavy elements Cu to Th. Astronomy & Astrophysics, 573, A27.
  • Grevesse, N., & Swings, J. P. (1969). Forbidden lines of Fe II in the solar photospheric spectrum. Astronomy & Astrophysics, 2, 28–44.
  • Hanasoge, S., Gizon, L., & Sreenivasan, K. R. (2016). Seismic sounding of convection in the Sun. Annual Review of Fluid Mechanics, 48, 191–217.
  • Harkins, W. D. (1917). The evolution of the elements and the stability of complex atoms. I. A new periodic system which shows a relation between the abundance of the elements and the structure of the nuclei of atoms. Journal of the American Chemical Society, 39, 856–879.
  • Heber, V. S., Baur, H., Bochsler, P., McKeegan, K. D., Neugebauer, M., Reisenfeld, D. B., . . . Wiens, R. C. (2012). Isotopic mass fractionation of solar wind: Evidence from fast and slow solar wind collected by the Genesis mission. Astrophysical Journal, 759, id121.
  • Heber, V. S., Burnett, D. S., Duprat, J., Guan, Y., Jurewicz, A. J. G., Marty, B., & McKeegan, K. D. (2013). Carbon, nitrogen and oxygen abundances in the bulk solar wind and calibration of absolute abundances. Lunar Planetary Science Conference 44, #2540. Houston, TX: Lunar and Planetary Institute.
  • Heber, V. S., McKeegan, K. D., Bochsler, P., Duprat, J., & Burnett, D. S. (2014). The elemental composition of solar wind with implications for fractionation processes during solar wind formation. Lunar Planetary Science Conference 45, #2117. Houston, TX: Lunar and Planetary Institute.
  • Heber, V. S., Wieler, R., Baur, H., Olinger, C., Friedmann, T. A., & Burnett, D. S. (2009). Noble gas composition of the solar wind as collected by the Genesis mission. Geochimica et Cosmochimica Acta, 73, 7414–7432.
  • Hill, F., Stark, P. B., Stebbins, R. T., Anderson, E. R., Antia, H. M., Brown, T. M., . . . Williams, W. E. (1996). The solar acoustic spectrum and eigenmode parameters. Science, 272, 1292–1295.
  • Holweger, H. (1979). Abundances of the elements in the Sun. In Les elements et leur isotopes dans l’univers: Communications presentees au XXIIe Colloque international d’astrophysique, tenu a Liege les 20, 21 et 22 Juin 1978 (pp. 117–138). Liège, Belgium: University of Liège, Institute of Astrophysics.
  • Holweger, H. (2001). Photospheric abundances, problems, updates, implications. In AIP Conference Proceedings Vol. 598 (pp. 23–30). College Park, MD: American Institute of Physics.
  • Holweger, H., & Müller, E. A. (1974). The photospheric barium spectrum, solar abundance and collision broadening of Ba II lines by hydrogen. Solar Physics, 39, 19–30.
  • Jurewicz, A. J. G., Burnett, D. S., Woolum, D. S., McKeegan, K. D., Heber, V., Guan, Y., . . .Hervig, R. (2011). Solar-wind Fe/Mg and a comparison with CI-chondrites. Lunar Planetary Science Conference 42, #1917. Houston, TX: Lunar and Planetary Institute.
  • Jurewicz, A. J. G., Olinger, C., Burnett, D. S., Rieck, K. D., & Woolum, D. S. (2019). SW Mg from Genesis: A new method of data reduction and implications. Lunar Planetary Science Conference 50, #2353. Houston, TX: Lunar and Planetary Institute.
  • Keil, K., & Fredriksson, K. (1964). The iron, magnesium, and calcium distribution in coexisting olivines and rhombic pyroxenes of chondrites. Journal of Geophysical Research, 69, 3487–3515.
  • Koeman-Shields, E. C., Huss, G. R., Ogliore, R. C., Jurewicz, A. J. G., Burnett, D. S., Nagashima, K., & Olinger, C. T. (2016). Hydrogen fluence calculated from Genesis collectors. Lunar Planetary Science Conference 47, #2800. Houston, TX: Lunar and Planetary Institute.
  • Laming, J. M. (2015). The FIP and inverse FIP effects in solar and stellar coronae. Living Reviews in Solar Physics, 12, id2.
  • Laming, J. M. (2017). The first ionization potential effect from the ponderomotive force: On the polarization and coronal origin of Alfvén waves. Astrophysical Journal, 844, id153.
  • Larimer, J. W., & Anders, E. (1967). Chemical fractionations in meteorites—II. Abundance patterns and their interpretation. Geochimica et Cosmochimica Acta, 31, 1239–1270.
  • Lawler, J. E., Feigenson, T., Sneden, C., Cowan, J. J., & Nave, G. (2018). Transition probabilities of Co II weak lines to the ground and low metastable levels. Astrophysical Journal Supplement Series, 238, id7.
  • Lawler, J. E., Hala, F. N., Sneden, C., Nave, G., Wood, M. P., & Cowan, J. J. (2019). Transition probabilities of Sc I and Sc II and Scandium abundances in the Sun, Arcturus, and HD 84937. Astrophysical Journal Supplement Series, 241, id21.
  • Lawler, J. E., Sneden, C., & Cowan, J. J. (2015). Improved Co I log(gf) Values and Abundance Determinations in the Photospheres of the Sun and Metal-poor Star HD 84937. The Astrophysical Journal Supplement, 220, id13.
  • Lawler, J. E., Sneden, C., Nave, G., Den Hartog, E. A., Emrahoğlu, N., & Cowan, J. J. (2017). Improved Cr II log(gf) values and abundance determinations in the photospheres of the Sun and metal-poor Star HD 84937. Astrophysical Journal Supplement Series, 228, id10.
  • Lawler, J. E., Wood, M. P., Den Hartog, E. A., Feigenson, T., Sneden, C., & Cowan, J. J. (2014). Improved V I log(gf) values and abundance determinations in the photospheres of the Sun and metal-poor Star HD 84937. Astrophysical Journal Supplement Series, 215, id20.
  • Levin, B. Y., Kozlovskaia, S. V., & Starkova, A. G. (1956). Sredrii khimicheskii Yostav meteoritov (The average chemical composition of meteorites). Meteoritika, 14, 38–53.
  • Lind, K., Amarsi, A. M., Asplund, M., Barklem, P. S., Bautista, M., Bergemann, M., . . . Pereira, T. M. D. (2017). Non-LTE line formation of Fe in late-type stars—IV. Modelling of the solar centre-to-limb variation in 3D. Monthly Notices of the Royal Astronomical Society, 468, 4311–4322.
  • Lodders, K. (2003). Solar system abundances and condensation temperatures of the elements. Astrophysical Journal, 591, 1220–1247.
  • Lodders, K. (2008). The solar argon abundance. Astrophysical Journal, 674, 607–611.
  • Lodders, K. (2019). The chemical composition of the solar system. In A. Formicola, M. Junker, L. Gialanella, & G. Imbriani (Eds.), Nuclei in the cosmos XV (Springer Proceedings in Physics 219, pp. 165–170). Cham, Switzerland: Springer.
  • Lodders, K., & Fegley, B. (1998). The planetary scientist’s companion. Oxford, UK: Oxford University Press.
  • Lodders, K., & Fegley, B. (2011). Chemistry of the solar system. London, UK: Royal Society of Chemistry.
  • Lodders, K., Palme, H., & Gail, H. P. (2009). Abundances of the elements in the solar system. In J. E. Trümper (Ed.), Landolt-Börnstein, New Series (Vol. VI/4B, pp. 560–630). Berlin, Germany: Springer Verlag.
  • Maas, Z. G., Pilachowski, C. A., & Hinkle, K. (2016). Chlorine abundances in cool stars. Astrophysical Journal, 152, id196.
  • Mahaffy, P. R., Donahue, T.M., Atreya, S.K., Owen, T.C., Niemann, H.B. (1998). Galileo probe measurements of D/H and He/ He in Jupiter’s atmosphere. Space Science Reviews, 84, 251–263.
  • Maiorca, E., Uitenbroek, H., Uttenthaler, S., Randich, S., Busso, M., & Magrini, L. (2014). A new solar fluorine abundance and a fluorine determination in the two open clusters M67 and NGC 6404. Astrophysical Journal, 788, id149.
  • Marsch, E., von Steiger, R., & Bochsler, P. (1995). Element fractionation by diffusion in the solar chromosphere. Astronomy & Astrophysics, 301, 261–276.
  • Marty, B., Zimmermann, L., Burnard, P. G., Wieler, R., Heber, V. S., Burnett, D. S., . . . Bochsler, P. (2010). Nitrogen isotopes in the recent solar wind from the analysis of Genesis targets: Evidence for large scale isotope heterogeneity in the early solar system. Geochimica et Cosmochimica Acta, 74, 340–355.
  • Mason, B. (1960). The origin of meteorites. Journal of Geophysical Research, 65, 2965–2967.
  • Mason, B. (1962). Meteorites. New York, NY: John Wiley.
  • Mason, B. (Ed.). (1971). Handbook of elemental abundances in meteorites. New York, NY: Gordon & Breach.
  • Mason, B., & Wiik, H. B. (1964). The amphoterites and meteorites of similar composition. Geochimica et Cosmochimica Acta, 28, 533–538.
  • Mason, G. M., Mazur, J. E., & Dwyer, J. R. (1999). 3He Enhancements in Large Solar Energetic Particle Events. The Astrophysical Journal, 525, L133–L136.
  • Mason, G. M., Mazur, J. E., Dwyer, J. R., Jokipii, J.R., Gold, R. E., & Krimigis, S. M. (2004). Abundances of heavy and ultraheavy ions in 3He-rich solar flares. Astrophysical Journal, 606, 555–564.
  • McDonough, W. F., & Sun., S. S. (1995). The composition of the Earth. Chemical Geology, 120, 223–253.
  • McKeegan, K. D., Kallio, A. P. A., Heber, V. S., Jarzebinski, G., Mao, P. H., Coath, C. D., . . . Burnett, D. S. (2011). The oxygen isotopic composition of the Sun inferred from captured solar wind. Science, 332, 1528–2011.
  • Meija, J., Coplen, T. B., Berglund, M., Brand, W. A., de Bievre, P., Groning, M., . . . Prohaska, T. (2016). Isotopic compositions of the elements 2013 (IUPAC Technical Report). Pure Applied Chemistry, 88, 293–306.
  • Menzel, D. H. (1931). A study of the solar chromosphere. Publications of the Lick Observatory Vol. 17. San Jose, CA: Lick Observatory.
  • Merrill, G. P. (1909). The composition of stony meteorites compared with that of terrestrial igneous rocks, and considered with reference to their efficacy in world-making. American Journal of Science, 27, 469–474.
  • Merrill, G. P. (1915). Researches on the chemical and mineralogical composition of meteorites. Proceedings of the National Academy of Sciences, 1, 429–431.
  • Merrill, G. P. (1916). A report on researches on the chemical and mineralogical composition of meteorites, with especial reference to their minor constituents. Memoirs of the National Academy of Sciences, 14, 7–25.
  • Meshik, A., Hohenberg, C., Pravdivtseva, O., & Burnett, D. (2014). Heavy noble gases in solar wind delivered by Genesis mission. Geochimica et Cosmochimica Acta, 127, 326–347.
  • Meyer, J. P. (1985a). Solar-stellar outer atmospheres and energetic particles, and galactic cosmic rays. Astrophysical Journal Supplement Series, 57, 173–204.
  • Meyer, J. P. (1985b). The baseline composition of solar energetic particles. Astrophysical Journal Supplement Series, 57, 151–171.
  • Milligan, R. O., & Ireland, J. (2018). On the performance of multi-instrument solar flare observations during solar cycle 24. Solar Physics, 293, id18.
  • Müller, E. A. (1965). The composition of the solar atmosphere. In J. N. Xanthakis (Ed.), Solar physics: Proceedings of the NATO Advanced Study Institute on Solar Physics, held at Lagonissi, Athens, Greece, September 1965 (pp. 33–67). New York, NY: John Wiley.
  • Müller, E. A. (1968). The solar abundances. In L. H. Ahrens (Ed.), Origin and distribution of the elements (pp. 155–176). Oxford, UK: Pergamon Press.
  • Noddack, I., & Noddack, W. (1930). Die Häufigkeit der chemischen Elemente. Naturwissenschaften, 18, 757–764.
  • Noddack, I., & Noddack, W. (1934). Die geochemischen Verteilungskoeffizienten der Elemente. Svensk Kemisk Tidskrift, 46, 173–201.
  • Noddack, I., & Noddack, W. (1935). Die Häufigkeit en der seltenen Erden in Meteoriten. Zeitschrift für anorganische und allgemeine Chemie, 225, 337–364.
  • Nordlander, T., & Lind, K. (2017). Non-LTE aluminum abundances in late-type stars. Astronomy & Astrophysics, 607, A75.
  • Nussbaumer, H., & Swings, J. P. (1970). [Fe II] magnetic dipole transition probabilities and the problem of the solar iron abundance. Astronomy & Astrophysics, 1, 455–458.
  • Oddo, G. (1914). Die Molekularstruktur der radioaktiven Atome. Zeitschrift für anorganische und allgemeine Chemie, 87, 253–268.
  • O’Dwyer, B., Del Zanna, G., Mason, H. E., Sterling, A. C., Tripathi, D., & Young, P. R. (2011). Hinode extreme-ultraviolet imaging spectrometer observations of a limb active region. Astronomy & Astrophysics, 525, A137.
  • Ott, U. (2015). Planetary and pre-solar noble gases in meteorites. Chemie der Erde, Geochemistry, 74, 519–544.
  • Pagel, B. E. J. (1979). Solar abundances. A new table (October 1976). In L. H. Ahrens (Ed.), Origin and distribution of the elements (pp. 79–80). Oxford, UK: Pergamon Press.
  • Palme, H. (1988). Chemical abundances in meteorites. In G. Klare (Ed.), Reviews in modern astronomy (pp. 28–51). Berlin, Germany: Springer.
  • Palme, H., & Beer, H. (1993). Abundances of the elements in the solar system. In H. H. Voigt (Ed.), Landolt-Bornstein, Group VI (pp. 196–221). Astronomy and Astrophysics Vol. 3, Extension and Supplement to Vol. 2, Subvolume A. Berlin, Germany: Springer Verlag.
  • Palme, H., & Jones, A. (2003). solar system abundances of the elements, meteorites, comets and planets. In A. M. Davis (Ed.), Treatise on geochemistry (Vol. 1, pp. 41–61). Amsterdam, The Netherlands: Elsevier.
  • Palme, H., Lodders, K., & Jones, A. (2014). solar system abundances of the elements. In A. M. Davis (Ed.), Treatise on geochemistry (2nd ed., Vol. 2, pp. 15–36). Amsterdam, The Netherlands: Elsevier.
  • Palme, H., Suess, H. E., & Zeh, H. D. (1981). Abundances of the elements in the solar system. In K.-H. Hellwege (Ed.), Landolt-Bornstein, Group VI (pp. 257–272). Astronomy and Astrophysics Vol. 2. Extension and Supplement I. Subvolume A. Berlin, Germany: Springer.
  • Parker, E. N. (1958). Dynamics of the interplanetary gas and magnetic fields. Astrophysical Journal, 128, 664–676.
  • Payne, C. H. (1925a). Astrophysical data bearing in the relative abundance of the elements. Proceedings of the National Academy of Sciences, 11, 192–198.
  • Payne, C. H. (1925b). Stellar atmospheres: A contribution to the observational study of high temperature in the reversing layers of stars (Doctoral dissertation). Harvard University, Cambridge, MA.
  • Pepin, R. O., Schlutter, D. J., Becker, R. H., & Reisenfeld, D. B. (2012). Helium, neon, and argon composition of the solar wind as recorded in gold and other Genesis collector materials. Geochimica et Cosmochimica Acta, 89, 62–80.
  • Pfeiffer, B., Venkataramaniahc, K., Czoka, U., & Scheidenberger, C. (2014). Atomic mass compilation 2012. Atomic Data Nuclear Tables, 100, 403–535.
  • Piersanti, L., Straniero, O., & Cristallo, S. (2007). A method to derive the absolute composition of the Sun, the solar system, and the stars. Astronomy & Astrophysics, 462, 1051–1062.
  • Pilleri, P., Reisenfeld, D. B., Zurbuchen, T. H., Lepri, S. T., Shearer, P., Gilbert, J. A., von Steiger, R., & Wiens, R. C. (2015). Variations in solar wind fractionation as seen by ACE/SWICS and the implications for GENESIS mission results. Astrophysical Journal, 812, id1.
  • Pillinger, C. T. (1979). Solar-wind exposure effects in the lunar soil. Reports on Progress in Physics, 42, 897–961.
  • Plaskett, H. H. (1922). Spectra of three O type stars. Publication Dominion Astronomical Observatory, 1, 325–384.
  • Pottasch, S. R. (1964a). A comparison of the chemical composition of the solar atmosphere with meteorites. Annales d’Astrophysique, 27, 163–169.
  • Pottasch, S. R. (1964b). On the chemical composition of the solar corona. Monthly Notices of the Royal Astronomical Society, 128, 73–91.
  • Pottasch, S. R. (1964c). On the interpretation of the solar ultraviolet emission line spectrum. Space Science Reviews, 3, 816–855.
  • Pottasch, S. R. (1967). The inclusion of dielectronic recombination processes in the interpretation of the solar ultraviolet spectrum. Bulletin of the Astronomical Institute of the Netherlands, 19, 113–124.
  • Prior, G. T. (1916). On the genetic relationships and classification of meteorites. Mineralogical Magazine, 18, 26–44.
  • Prior, G. T. (1920). The classification of meteorites. Mineralogical Magazine, 19, 51–63.
  • Rankama, K. (1956). Isotope geology. London, UK: Pergamon Press.
  • Reames, D. V. (2018). Abundances, ionization states, temperatures, and FIP in solar energetic particles. Space Science Reviews, 214, id61.
  • Reed, G. W., Kigoshi, K., & Turekevitch, A. (1960). Determinations of concentrations of heavy elements in meteorites by activation analysis. Geochimica et Cosmochimica Acta, 20, 122–140.
  • Reisenfeld, D. B., Wiens, R. C., Barraclough, B. L., Steinberg, J. T., Neugebauer, M., Raines, J., & Zurbuchen, T. H. (2013). Solar wind conditions and composition during the Genesis mission as measured by in situ spacecraft. Space Science Reviews, 175, 125–164.
  • Rieck, K., Jurewicz, A. J. G., Burnett, D. S., Hervig, R. L., Williams, P., & Guan, Y. (2016). Bulk solar wind Na and K measured in Genesis collectors. Lunar Planetary Science Conference 47, #2922. Houston, TX: Lunar and Planetary Institute.
  • Rieck, K. C. Ogliore, R. C., Jurewicz, A. J. G., Burnett, D. S., Guan, Y., & Wiens, R. C. (2019). Measuring solar wind C and O abundances in Genesis regime collectors using SIMS ion imaging depth profiling. Lunar and Planetary Science Conference, 50, 2944.
  • Ringwood, A. (1961). Chemical and genetic relationships among meteorites. Geochimica et Cosmochimica Acta, 24, 159–197.
  • Rose, G. (1863–1864). Beschreibung und Eintheilung der Meteoriten. Abhandlungen der Königlichen Akademie der Wissenschaften in Berlin, 23–161.
  • Ross, J. E., & Aller, L. H. (1976). The chemical composition of the Sun. Science, 191, 1223–1229.
  • Russell, H. N. (1914). The solar spectrum and the Earth’s crust. Science, 39, 791–794.
  • Russell, H. N. (1929). On the composition of the Sun’s atmosphere. Astrophysical Journal, 70, 11–28.
  • Saslawsky, I. I. (1933). Die Zusammensetzung der Meteorite und des Erdballes und die Kontraktionsgroesse des Erdballes. Tschermak’s Mineralogische Petrologische Mitteilungen, 43, 144–155.
  • Schmelz, J. T., Reames, D. V., von Steiger, R., & Basu, S. (2012). Composition of the solar corona, solar wind, and solar energetic particles. Astrophysical Journal, 755, id33.
  • Schou, J., Christensen-Dalsgaard, J., Howe, R., Larsen, R. M., Thompson, M. J., & Toomre, J. (1998). In S. G. Korzennik & A. Wilson (Eds.), Proceedings of SOHO 6/GONG 98 Workshop: Structure and Dynamics of the Interior of the Sun and Sun-like Stars (Vol. 2, pp. 845). ESA SP-418. Noordwijk, The Netherlands: European Space Agency.
  • Scott, E. R. D., & Krot, S. (2014). Chondrites and their components. In A. M. Davis (Ed.), Treatise on geochemistry (2nd ed., Vol. 2, pp. 65–137). Amsterdam, The Netherlands: Elsevier.
  • Scott, P., Asplund, N., Grevesse, N., Bergemann, M., & Sauval, A. J. (2015). The elemental composition of the Sun. II. The iron group elements Sc to Ni. Astronomy & Astrophysics, 573, A26.
  • Scott, P., Grevesse, N., Asplund, M., Sauval, A. J., Lind, K., Takeda, Y., . . . Hayek, W. (2015). The elemental composition of the Sun. I. The intermediate mass elements Na to Ca. Astronomy & Astrophysics, 573, idA25.
  • Serenelli, A. M., & Basu, S. (2010). Determining the initial helium abundance of the Sun. Astrophysical Journal, 719, 865–872.
  • Serenelli, A. M., Haxton, W. C., & Pena-Gray, C. (2011). Solar models with accretion. I. Application to the solar abundance problem. Astrophysical Journal, 743, id24.
  • Sneden, C., Cowan, J. J., & Gallino, R. (2008). Neutron-capture elements in the early galaxy. Annual Review of Astronomy and Astrophysics, 46, 241–288.
  • Sneden, C., Cowan, J. J., Kobayashi, C., Pignatari, M., Lawler, J. E., Den Hartog, E. A., & Wood, M. P. (2016). Iron-group abundances in the metal-poor main-sequence turnoff star HD84937. Astrophysical Journal, 817, id53.
  • Sneden, C., Lawler, J. E., Cowan, J. J., Ivans, I. I., & Den Hartog, E. A. (2009). New rare earth element abundance distributions for the Sun and five r-process-rich very metal-poor stars. Astrophysical Journal Supplemental Series, 182, 80–96.
  • Strömgren, B. (1940). On the chemical composition of the solar atmosphere. In K. Lundmark (Ed.), Festschrift für Elis Strömgren (pp. 218–257). Cophenhagen, Denmark: Einar Munksguard.
  • Suess, H. E. (1947a). Über kosmische Kernhäufigkeiten. I. Mitteilung, Einige Häufigkeitsregeln und ihre Anwendung bei der Abschätzung der Häufigkeitswerte für die mittelschweren und schweren Elemente. Zeitschrift für Naturforschung, 2a, 311–321.
  • Suess, H. E. (1947b). Über kosmische Kernhäufigkeiten II. Mitteilung, Einzelheiten in der Häufigkeitsverteilung der mittelschweren und schweren Kerne. Zeitschrift für Naturforschung, 2a, 604–608.
  • Suess, H. E. (1949). Die kosmische Häufigkeit der chemischen Elemente. Experimentia, 7, 266–270.
  • Suess, H. E., & Urey H. C. (1956). Abundances of the elements. Reviews in Modern Physics, 28, 53–74.
  • Suess, H. E., & Zeh, H. D. (1973). The abundances of the heavy elements. Astrophysics and Space Science, 23, 173–187.
  • Sylwester, B., Phillips, K. J. H., Sylwester, J., & Kępa, A. (2015). Resik solar x-ray flare element abundances on a non-isothermal assumption. Astrophysical Journal, 805, id 49.
  • Tousey, R. (1967). Some results of twenty years of extreme ultraviolet solar research. Astrophysical Journal, 149, 239–252.
  • Trieloff, M. (2018). Noble gases. In Oxford Research Encyclopedia of Planetary Science. Oxford University Press.
  • Turck-Chieze, S., Couvidat, S., Piau, L., Ferguson, J., Lambert, P., Ballot, J., . . . Nghiem, P. (2004). Surprising Sun: A new step towards a complete picture? Physical Review Letters, 93, id 211102.
  • Turcotte, S., Richter, J., Michaud, G., Iglesia, C. A., & Rogers, F. J. (1998). Consistent solar evolution model including diffusion and radiative acceleration effects. Astrophysical Journal, 504, 539–558.
  • Turcotte, S., & Wimmer-Schweingruber, R. F. (2002). Possible in situ tests of the evolution of elemental and isotopic abundances in the solar convection zone. Journal of Geophysical Research, 107(A12), 1442.
  • Unsöld, A. (1945). Quantitative Spektralanalyse der Sonnenatmosphäre. Nachrichten Gesellschaft Wissenschaften, Göttingen, Germany: Mathematisch Physikalsiche Klasse, 69–70.
  • Unsöld, A. (1948). Quantitative Analyse der Sonnenatmosphäre. Zeitschrift für Astrophysik, 24, 306–329, erratum Zeitschrift für Astrophysik, 25, 352.
  • Urey, H. C. (1952). The abundances of the elements. Physical Review, 88, 248–252, erratum [1953], Physical Review, 90, 1131.
  • Urey, H.C. (1967). The abundance of the elements with special reference to the problem of the iron abundance. Quarterly Journal Royal Astronomical Society, 8, 23–47.
  • Urey, H. C. (1972). Abundances of the elements. Annals New York Academy of Sciences, 194, 35–44.
  • Urey, H. C., & Craig, H. (1953). The composition of the stone meteorites and the origin of the meteorites. Geochimica et Cosmochimica Acta, 4, 36–82.
  • Urey, H. C., & Edwards, G. (1956). Determination of alkali metals in meteorites by a distillation process. Geochimica et Cosmochimica Acta, 7, 154–168.
  • van Schmus, W. R., & Wood, J. A. (1967). A chemical-petrologic classification for the chondritic meteorites. Geochimica et Cosmochimica Acta, 31, 747–754.
  • Vauclair, S. (1998). Element settling in the solar interior. Space Science Reviews, 85, 71–78.
  • Veck, N. J., & Parkinson, J. H. (1981). Solar abundances from X-ray flare observations. Monthly Notices of the Royal Astronomical Society, 197, 41–55.
  • Vinyoles, N., Serenelli, A. M., Villante, F. L., Basu, S., Berstron, J., Gonzales-Garcia, M. C., . . . Song, N. (2017). A new generation of standard solar models. Astrophysical Journal, 835, id202.
  • Vitas, N., Vince, I., Lugaro, M., Andriyenko, O., Gosic, M., & Rutten, R. J. (2008). On the solar abundance of indium. Monthly Notices of the Royal Astronomical Society, 384, 370–375.
  • Vogel, N., Heber, Veronika, V. S., Baur, H., Burnett, D. S., & Wieler, R. (2011). Argon, krypton, and xenon in the bulk solar wind as collected by the Genesis mission. Geochimica et Cosmochimica Acta, 75, 3057–3071.
  • von Steiger, R., & Schwadron, N. A. (2000). Solar wind composition. In R. Ramaty & N. Mardzhavidz (Eds.), High energy solar physics: Anticipating HESSI (pp. 54–63). ASP Conference Series 206. San Francisco, CA: Astronomical Society of the Pacific.
  • von Steiger, R., Schwadron, N. A., Fisk, L. A., Geiss, J., Gloeckler, G., Hefti, S., . . . Zurbuchen, T. H. (2000). Composition of quasi-stationary solar wind flows from Ulysses/Solar Wind Ion Composition Spectrometer. Journal of Geophysical Research, 105(A12), 27217–27238.
  • von Steiger, R., & Zurbuchen, T. H. (2016). Solar metallicity derived from in situ solar wind composition. Astrophysical Journal, 816, id13.
  • Wahl, W. A. (1910). Beitraege zur Chemie der Meteoriten (Contributions about the chemistry of meteorites). Zeitschrift für anorganische Chemie, 69, 52–96.
  • Wasson, J. T. (1985). Meteorites—Their record of the early Solar-System. New York, NY: W. H. Freeman.
  • Wasson, J. T., & Kallemeyn, G. W. (1988). Compositions of chondrites. Philosophical Transactions of the Royal Society A, 325, 535–544.
  • Westphal, A. J., Huss, G. R., Koeman-Shields, E., Welten, K., & Westphal, L. K. T. (2019). Abundance of manganese in the solar wind using dual-implant calibration. Lunar Planetary Science Conference 50, #2016. Houston, TX: Lunar and Planetary Institute.
  • Widing, K. G., & Feldman, U. (2001). On the rate of abundance modifications versus time in active region plasmas. Astrophysical Journal, 555, 426–434.
  • Wieler, R. (2016). Do lunar and meteoritic archives record temporal variations in the composition of solar wind noble gases and nitrogen? A reassessment in the light of Genesis data. Chemie der Erde, Geochemistry, 76, 463–480.
  • Wiik, H. B. (1956). The chemical composition of some stony meteorites. Geochimica et Cosmochimica Acta, 9, 279–289.
  • Wildt, R. (1939). Negative ions of hydrogen and the opacity of stellar atmospheres. Astrophysical Journal, 90, 611–620.
  • Withbroe, G. L. (1971). The chemical composition of the photosphere and corona. In K. B. Gebbie (Ed.), The Menzel Symposium on Solar Physics, Atomic Spectra, and Gaseous Nebulae (pp. 127–162). National Bureau of Standards Special Publication No. 353. Washington, DC: U.S. Department of Commerce, National Bureau of Standards.
  • Withbroe, G. L. (1975). The analysis of XUV emission lines. Solar Physics, 45, 301–317.
  • Woolley, R. D. V. R., & Allen, C. W. (1948). The coronal emission spectrum. Monthly Notices of the Royal Astronomical Society, 108, 292–305.
  • Yang, W. (2016). Solar models with new low metal abundances. Astrophysical Journal, 821, 108.
  • Yang, W. (2019). Rotating solar models with low metal abundances as good as those with high metal abundances. Astrophysical Journal, 873, id18.
  • Yavnel, A. A. (1968). On the equilibrium of the composition of the iron-magnesium silicates in ordinary chondrites. Meteoritika, 28, 19–29.
  • Yavnel, A. A. (1970). On the degree of equilibration reflected in the ferromagnesian silicates in ordinary chondrites, Meteoritics, 5, 153–167.
  • Young, P. R. (2018). Element abundance ratios in the quiet Sun transition region. Astrophysical Journal, 855, id15.
  • Zhang, Q.-S., Li, Y., & Christensen-Dalsgaard, J. (2019). Solar models with convective overshoot, the solar-wind mass loss, and a PMS disk accretion: Helioseismic quantities, Li depletion and neutrino fluxes. Astrophysical Journal, 881, 103.
  • Zhao, G., Mashonkina, L., Yan, H. L., Alexeeva, S., Kobayashi, C., Pakhomov, Y., . . . Zhai, M. (2016). Systematic non-LTE study of the −2.6 ≤ [Fe/H] ≤ 0.2 F and G dwarfs in the solar neighborhood. II. Abundance patterns from Li to Eu. Astrophysical Journal, 833, 225.
  • Zurbuchen, T. H., Weberg, M., von Steiger, R., Mewaldt, R. A., Lepri, S. T., & Antiochos, S. K. (2016). Composition of coronal mass ejections. Astrophysical Journal, 826, id10.