This element was known in antiquity. In Latin iron is called ferrum, therefore the abbreviation Fe.
FeI 7.9 eV, FeII 16.2 eV, FeIII 30.6 eV, FeIV 54.8 eV, FeV 75.5 eV, FeVI 100 eV, FeVII 128.3 eV, FeVIII 151.1 eV, FeIX 235 eV, FeX 262 eV.
Fe is an element whose lines dominate the spectra of all stars of type F and later. In the solar spectrum, iron lines account for about 30% of the lines, a percentage not surpassed by any other element. Because of this, a large part of the line blocking is due to Fe and therefore Fe was taken to represent all elements other than H and He (see also Part Two, section 2.1).
Absorption lines of FeI
FeI lines (for instance 4045, 5216 and 5269) appear in late B-type stars and grow rapidly toward later types. No luminosity effect is observed.
Among the most intense infrared lines one has 10216 (1247). In the sun, W(10216) = 0.094.
High-excitation-potential Fe I lines (about 7 eV) were detected in the sun and in a late type giant at 2550-2500 cm-1 Johansson et al. 1991).
Emission lines of FeI
FeI emissions are found in the ultraviolet spectrum of K giants and M-type supergiants (van der Hucht et al. 1979).
In T Tau stars several multiples (for instance M.2) are seen in emission Joy 1945). In addition the lines 4063 and 4132 (both of M.43) are often unexpectedly strong. This is due to fluoresence excited either by CaII or by H epsilon (Willson 1975).
In long-period variables many lines of Fe I are seen in emission, usually around maximum light; they disappear about at the minimum Joy 1954).
Particularly strong are the 4202 and 4308 lines, both of M.42 (Querci 1986).
Several lines of Fe I are seen in emission in the superaova 1987A (Arnett et al. 1989).
Absorption lines of FeII
The 4233 line is present up to M 2 III (Davis 1947).
FeII (see for example 4233) lines appear in mid-B-type stars and strengthen toward A-type stars. They have a maximum for late F and decrease for G type. A positive luminosity effect is present, which is much used for spectral classification.
The Fe II lines 6149 and 6147 (M.74) are very sensitive to the Zeeman effect, but in different ways, because of their different splitting pattern and the operation of the Paschen-Back effect (Mathys 1990). They have been used to measure the magnetic field in Ap stars (Mathys and Lanz 1992).
Emission lines of Fell
FeII emissions are ubiquitous and one finds them whenever a low-density medium is heated by shock waves, expanding atmospheres or radiation from hot spots or a binary companion. They are thus present in extended envelopes and in stellar chromospheres. Many details on FeII and [Fe II] emissions can be found in the colloquium Physics of Formation of FeII Lines Outside LTE edited by Viotti and Friedjung (1988). FeII is seen in emission in the infrared in Oe stars (Andrillat et al. (1982). Fe II and (Fe II) emissions are present in variable OB stars (S Dor type) and P Cyg stars (Zickgraf 1988). (For the latter see also Stahl et al. (1991).) Often both types of objects are listed together under the title of luminous blue variables. For a summary of their properties, see Wolf (1993).
Fe II lines (M.42) are seen in emission in early Be stars (Slettebak et al. 1992), but an inspection of the infrared region shows that Fe II emissions (for instance 9997) are present in all Be stars, with an intensity depending upon the spectral type. This is probably due to fluoresence (Andrillat et al. 1990). W(9997) can reach up to 1 ångström. [Fe II] lines in emission are characteristic of B[e] stars (Andrillat and Swings 1976). These stars also show permitted FeII lines in emission. FeII emissions are also seen in Herbig Ae-Be stars (Talavera 1988).
Numerous (faint) emission lines of FeII are seen in the spectrum of the chromosphere of late type stars of types K and M (van der Hucht et al. 1979, Linsky et al. 1982, Judge and Jordan 1991). In one K giant these lines are responsible for more than 80% of all emission lines (Carpenter and Wing 1985).
In symbiotic stars one observes high-excitation ultraviolet fluoresence Fell emissions (1360, 1776, 1869, 1881, 1884 and 1975) produced by the strong O VI, CIV and HI lines (Feibelman et al. 1991). FeII lines are also present. In the infrared the 9997 line is usually seen in emission (Baratta et al. 1991).
T Tau stars often exhibit in emission the lines of M.27, 28, 37 and 38, besides those of the chromospheric spectrum (Sun etal. 1985, Jordan 1988). In addition some [Fe II] lines are also seen in emission (Joy 1945).
In semiregular variables Fe II is always seen in emission, with time variations of the line strength (Querci and Querci 1989).
In Mira-type variables, Fe II emissions (for instance M.27 and 28) are present over a large part of the cycle, except for a few weeks after minimum (Joy 1954). Some forbidden FeII lines (M.6F and 7F) are visible around the minimum in long-period variables of type S. Many FeII lines appear in emission around maximum light: 4233,(M.27), 4583(37), 4924 and 5018(42) (Merrill 1952).
FeII emissions are also found in R CrB stars and in binaries with a hot companion, like cataclysmic variables, dwarf novae, symbiotic stars,VV Cep stars (Viorti and Friedjung 1988) and in novae (Cassatella and Gonzalez Riestra 1988) (for the novae see also below) and in supernovea 1987A (Arnett et al. 1989). [FeII] emission lines are also present in VV Cep stars (Rossi et al. 1992) and in supernovae 1987A (Jennings et al. 1993).
Absorption lines of FeIII
|Source: Some data are from Kilian and Nissen (1989).|
FeIII (see for example 4164) appears in late O- and early B-type stars. For dwarfs the maximum lies around B 0. A positive luminosity effect exists. Fe III is prominent in the ultraviolet of early type stars, as is the case for FeII.
Emission lines of FEIII
Lines of Fe III appear in emission in the ultraviolet spectrum of K and M giants and supergiants (van der Hucht et al. 1979). [FeIII] is seen often in B[e] stars (Andrillat and Swings 1976) and [FeIII] lines 4658 and 5270 have been observed in some novae (Swings 1952).
Absorption lines of FeIV, Fe V and Fe VI
Many weak lines of FeIV, FeV and FeVI appear in the ultraviolet spectrum of one O-type subdwarf (Bruhweiler et al. 1981), the FeVI lines being the weakest. Nemry et al. (1991) found weak FeIV and FeV lines also in one WR star( 0.1).
Emission lines of Fel V, Fe V and Fe Vl
Emissions of forbidden FeV(3895, 3891) and FeVI (3664) have been observed in some novae and some symbiotic stars (Merrill 1948, Swings 1952).
Lines of Fe VII
Joy and Swings (1945) reported FeVII in one recurrent nova. The 6087 line of [FeVII] is always seen in emission in symbiotic stars. It should, however, be remarked that often one sees [FeII] and [FeVII] without the intermediate species.
Wallerstein et al. (1991) also report observation of the 3586 line of [FeVII] in emission in several symbiotic stars.
Forbidden lines of higher ionization stages
Forbidden lines of FeX, XI, XII, XIII, XIV and XV (7060) with ionization potentials between 233 and 390 eV are seen in the spectrum of the solar corona (Zirin 1988). (See also Part Two, section 3.2.)
[FeX] and [FeXIV] in emission have been observed in some recurrent novae (Joy and Swings 1945, Swings 1952).
Behavior in non-normal stars
Fe is weak in some of the extreme helium stars Jeffery and Heber 1993).
Fe appears enhanced in Bp stars of Si type. For instance FeII 4233 is intensified by factors up to three, but there is a continuous transition to normal stars (Didelon 1986). Fe, on the other hand, may be very weak in some Hg-Mn stars (for instance 53 Tau), where even strong ultraviolet multiplets are missing. Nevertheless, in other stars of the Hg-Mn subgroup Fe lines are strong.
In Ap stars of the Cr-Eu-Sr subgroup there also exist cases of stars in which iron lines are strong (Cowley 1981). In some stars of this subgroup one finds iron simultaneously in three ionization stages, namely FeI, FeII and FeIII (Jaschek and Lopez Garcia 1966).
Fe is slightly strengthened in Am stars (Burkhart and Coupry 1991).
Fe lines are weakened in lambda Boo stars. Typically the W values are smaller by factors of about two than in normal stars of the same temperature (Venn and Lambert 1990).
Fe lines are weakened in F-type HB stars by factors up to five in W values (Adelman and Hill 1987).
Fe I lines are present in the spectra of some DZ degenerates (Sion et al. 1990), with W(3820 < 5.5 Å).
Fe is also underabundant in metal-weak stars. It should be remarked that generally authors speak of `metal abundances', rather than of iron abundances.
Fe is slightly weak in subgiant CH stars (Luck and Bond 1982).
Fe is underabundant in globular cluster stars, varying between -0.8 dex for `metal-rich' globular clusters, to about -1.3 dex in `intermediate clusters' and -2.3 dex or less in `metal-poor' clusters.
The Fe/H ratio seems to be normal (solar) in the SMC stars. In the LMC stars real variations in Fe/H may exist between different stars (Luck and Lambert 1992).
Fe is slightly underabundant in Ba stars (Lambert 1985) and is probably normal in S-type stars (Lambert 1989).
Forbidden iron lines are visible in novae, especially in the 'nebular stage' phase. Their behavior varies from star to star. In some objects one sees lines from [Fe II] up to [Fe VII] whereas other stars have shown lines from [Fe XI] to [Fe VIII], with [FeXI] appearing before [FeVIII] (Warner 1989).
Forbidden iron lines are often seen in symbiotic stars - for instance from Fe II, III, IV, VI and VII (for instance Barba et al. (1992), Freitas Pacheco and Costa (1992)).
In supernovae, absorption lines of Fe II appear in all subtypes (Ia, Ib and II). Emission lines of both [FeII] and [FeIII] are present in the nebular phase of supernovae of type Ia (Branch 1990).
Fe occurs in four stable isotopes, one long-lived and five short-lived ones. The stable isotopes are Fe 54, 56, 57 and 58, which occur in the solar system in 6%, 92%, 2% and 0.3% abundances. The long-lived isotope is Fe60 with a half life of 3 × 105 years.
Fe54 is produced by explosive nucleosynthesis, Fe56 and Fe57 also by this process or by a statistical equilibrium process; Fe58 can be produced by either helium burning or a nuclear statistical equilibrium process and also by carbon burning.
Published in "The Behavior of Chemical Elements in Stars", Carlos Jaschek and Mercedes Jaschek, 1995, Cambridge University Press.