This element was first isolated by H. Davy in London in 1808. The name comes from the Greek city Magnesia, where a mineral known to the ancients was found.

Ionization energies
MgI 7.6 eV, MgII 15.0 eV, MgIII 80.0 eV, MgIV 109.2 eV, MgV 141.3 eV, MgVI 186.5 eV, MgVII 225 eV.

Absorption lines of MgI

Table 1: Equivalent widths of MgI

  5172(2)     5711(8)    

Group V III Ib V III Ib

B9 0.12          
A0 0.15          
A7 0.25          
FO 0.41          
F4 0.48          
F5 0.485   0.42 0.069   0.064
F6 0.53          
F8 0.64   0.615     0.10
G0       0.12   0.15
G1 0.78   0.534      
G2 1.09,1.30   0.640     0.16
S 1.259          
G5 1.23   0.770 0.11   0.18
G8   1.29       0.16
K0 2.08 1.95,1.73   0.17 0.19 0.14
K2 2.08 1.45       0.22
K3         0.17 0.27
K5       0.20   0.22
M0         0.15  
M2           0.18

Mg I (see for instance 5172) appears in A-type stars and increases toward later types. A negative luminosity effect appears in the 5172 line from late F-type on. In contrast, a positive luminosity effect is present in the 5711 line.

The 5172 feature has such a large equivalent width that it can also be measured photoelectrically, see for instance Guinan and Smith (1984).

Fanelli et al. (1990) have analyzed the behavior of 2852 (UV M.1). The line appears around A0, increases slowly toward F7 and then rises rapidly to a maximum at K3. The line has a negative luminosity effect, like that of the line at 5172, after F7.

Figure 30

Emission lines of MgI
MgI emissions appear in the chromospheric spectrum of K and M giants and supergiants (van der Hucht et al. 1979).

Some lines of MgI appear in emission in T Tau stars (Joy 1945, Appenzeller et al. 1980).

In long-period variables of type S many Mg I lines appear in emission around the maximum - for instance 3829, 3832, 3838(3) and 5167, 5172 and 5183(2). Near the minimum the line 4571(1) is in emission in all long-period variables (Merrill 1952).

The 2852(1) line appears in emission in some M-type supergiants. Fix and Cobb (1987) found MgI emission lines in the 1.6 µm region of an F-type supergiant that is an OH maser.

Absorption lines of MgII

Table 2: Equivalent widths of MgII

  4481(4)     2800    

Group V III Ia V III Ib

O7 0.044          
O9 0.074          
B0 0.104   0.150      
B1 0.133   0.16 0.9    
B3 0.19   0.404 1.3    
B5 0.225 0.29 0.444      
B6 0.343 0.395        
B7       2.4    
B8     0.43 4.0    
B9 0.39          
A0 0.4   0.59(Ia)      
A1 0.41          
A2 0.43,0.38   0.676,0.913(0)      
A3     0.50(Ib),0.573(0) 7.0    
A5         8.9  
A7 0.50          
F0 0.31   0.776      
F1           6.9
F2     0.468(Ib) 9.9    
F4 0.36          
F5 0.33,0.24   0.48(Ib) 8.4    
F6 0.35          
F8 0.29   0.44(Ib)      
GO 0.20          
G1 0.28          
G2 0.125          
S 0.120          
G5 0.22          

Note: The 2800 feature is a blend of 2790(3), 2795(1), 2798(3) and 2802(1).

Source: The 2800 line data are from Kondo et al. (1977). This line's behavior was also studied by Fanelli et al. (1990), but they do not provide equivalent widths.

MgII (see for instance 4481) appears in B-type stars, has a maximum at late A types and a positive luminosity effect. The strong ground level transitions at 2800 (2795 and 2802), which are analogous to the H and K lines of CaII, show a similar behavior, but the matter should be investigated in more detail with more stars. For the latter see also Gurzadyan (1975).

Figure 31

Mg II emissions

Emissions in the MgII resonance doublet 2795 and 2803 are similar in appearance to those of the CaII doublet, showing a centrally reversed emission core within the broad absorption line. (For illustrations of the line profiles see for instance Basri and Linsky (1979).) In general MgII emissions appear in the same type of stars as CaII emissions, namely late type stars of all luminosity classes from F-type onwards. They become very frequent from type K 3 onwards. In general the MgII emissions are the strongest ones in the near ultraviolet region. Their presence signals the existence of a stellar chromosphere. (See Part Two, section 3.1.) Their presence and strength are thus not correlated directly with metal weakness (Dupree et al. 1990) as was thought for some time. Several surveys of these emissions have been made, among others by Stencel et al. (1980) and Doherty (1985).

The behavior of the 2800 feature is similar to that of Ca II and one can derive a relation between the width of the emission feature (d) and the absolute magnitude of the star (Kondo et al. 1977). Elgoroy (1988) gives

Mv = - 19.58 log d + 41.36

This relation is analogous to the Wilson-Bappu relation for CaII.

The MgII lines 7877 and 7896 of M 8 appear in emission in some early type peculiar objects, like Be, B[e] and Herbig Ae-Be stars Jaschek et al. 1993). In early Be stars one sees also 9217-9243(1) in emission Johnson 1977).

MgII in emission is often observed in T Tau stars (Appenzeller et al.1980) and always in semiregular variables (Querci and Querci 1989). Profiles of the MgII emissions have also been used to map the bright plages (bright emission regions) in RS CVn stars. For a summary see Rodono (1992).

Behavior in non-normal stars

Mg is overabundant by 0.6 dex in the spectrum of at least one extreme He star (Schoenberner and Drilling 1984).

The 4481(4) line of MgII is weakened in shell spectra (Struve and Wurm 1938).

Bidelman (1966) called attention to one early Ap star characterized by strong SiII and MgII lines. Jaschek and Jaschek (1974) later listed a small subgroup of stars that share this characteristic and can be called Mg-strong stars. For a detailed study of one star of the group see Naftilan (1977).

MgII tends to be weak in Ap stars of the Cr-Eu-Sr subgroup (Adelman 1973b). Typically W(4481) = 0.30 (Sadakane 1976).

Mg tends also to be weak in Am stars. Smith (1973, 1974) found that W values in Am stars are about 70% those of normal stars.

MgII lines are generally weak in lambda Boo stars (Holweger and Stuerenburg 1991). Typically W(4481) = 0.100 for A0V. Abt (1984) has used the Mg weakness as a criterion to select candidates for new lambda Boo stars.

Mg lines are weakened in F-type HB stars with regard to normal stars of the same colors, by factors up to ten in W (Adelman and Hill 1987). MgI(3832, 3838) is seen in DZ degenerates (Sion et al. 1990), with W up to 8Å.

Mg seems to be slightly overabundant with respect to iron in metal-weak stars by factors of the order of two (Gratton and Sneden 1987, Magain 1989, Francois 1991). This is also true for the most metal-deficient stars, with Fe/H about -4 dex (Molaro and Bonifacio 1990). In globular cluster stars Wheeler et al. (1989) found that Mg is underabundant with respect to Fe, but Smith and Wirth (1991) remarked that the behavior of Mg is somewhat erratic, if stars of different CN strength are examined.

In novae, MgII appears both in absorption and in emission. Emissions are usually seen in the ultraviolet spectral region. Mg appears again at the nebular stage phase in the form of [MgV] and [MgVII] (Payne-Gaposchkin 1957, Warner 1989).

Strong Mg II lines appear in the spectra of supernovae of types Ia and II. Strong emissions of [MgI] are present in the nebular stage of supernovae of type Ib (Branch 1990).

Mg Isotopes
Magnesium has three stable isotopes, Mg24, 25 and 26, which occur in the solar system with respectively 79%, 10% and 11% abundance. There also exist five unstable isotopes.

Since the bands of Mg24H, Mg25H and Mg26H lie close together, they can be used for an analysis of the comparative abundance of the three Mg isotopes. Such an analysis was performed for two Ba stars by Tomkin and Lambert (1979) and the isotope ratios found are compatible with the solar values. The analysis was extended by Malaney and Lambert (1988) to more Ba stars, with the same result (i.e. solar ratios). Smith and Lambert (1988) investigated in a similar way the MS and S stars and also found no obvious anomalies in the isotopic ratio. McWilliam and Lambert (1988) concluded from an analysis of old disk stars that the isotope ratio is about solar. It is possible that, in metal-weak stars, the isotopes Mg25 and Mg26 diminish with respect to Mg24.

Mg 24 can be produced either by hot hydrogen burning or by explosive nucleosynthesis. Both Mg25 and Mg26 can be produced by neon burning, carbon burning or explosive nucleosynthesis.

Published in "The Behavior of Chemical Elements in Stars", Carlos Jaschek and Mercedes Jaschek, 1995, Cambridge University Press.