This element was discovered by D. Rutherford in 1772 in Edinburgh. Its name comes from nitron (nitre) and gene (which forms).
NI 14.5 eV, NII 29.6 eV, NIII 47.4 eV, NIV 77.5 eV, NV 97.9 eV, NVI 552 eV, NVII 667 eV.
Absorption lines of NI
Source: Values for 8683 from Roby and Lambert (1990) and Sadakane and Okyudo (1989). Values for 8711 are from Jaschek andJaschek (unpublished).
The most useful lines are those of the infrared region. NI appears around mid-B-type stars and is seen in A and F stars with a maximum in A-type stars (see for instance the line at 8683). A strong positive luminosity effect exists (see for instance the line at 8711). The resonance line of NI (UV 1) is at 1200.
Emission lines of NI
NI (M.1 and 8) are often seen in emission in early Be stars. The W values are less than 1 Å. In Be skell stars the lines are generally in absorption (Andrillat et al. 1990).
The ultraviolet lines 1134(2) and 1167-70(6) are seen in strong emission in the solar spectrum.
Absorption lines of NII
NII lines (see for instance the line at 4630) are seen in early type stars with a maximum at B 2. A strong positive luminosity effect exists. The same behavior is shared by the NII lines 5667, 5676 and 5711(3).
Emission lines of NII
The [NII] line is present in emission in P Cyg (Stahl et al. 1991). Swings (1973) observed many [N II] lines in the spectrum of one typical B[e] star. The [N II] lines 6548 and 6584 have been observed in emission in two peculiar G-type super-giants (Sheffer and Lambert 1992). The [NII] line 6584 is prominent in the post-nova phase of novae (Warner 1989) and in symbiotic stars (Freitas Pacheco and Costa 1992). In luminous blue variables several faint [NII] emissions have been observed (Wolf 1993).
The ultraviolet lines 916(2) and 1085(1) appear in emission in the solar spectrum (Feldman and Doschek 1991).
Absorption lines of NIII
NIII lines are seen in O- and early B-type stars, with a maximum at O 9. A strong positive luminosity effect exists. The ultraviolet line 1758 has a maximum around B 0 and disappears toward B 5 (Prinja 1990).
Emission lines of NIII
For the emissions of N III in O-type stars see the discussion of behavior in non-normal stars.
NIII appears in emission in some symbiotic stars. The ultraviolet lines of M.1, 989-91 appear in emission in the solar spectrum (Feldman and Doschek 1991).
Absorption lines of NIV
Source: The only supergiant is plotted as a cross in the figure.
NIV lines are seen in O -type stars. A strong luminosity effect exists. NIV 6381 reaches its maximum strength at O 4-O 5 in supergiants. The ultraviolet NIV line 1718(7) has a maximum at O 4 (Heck et al. 1984).
Absorption lines of NV
NV (resonance line at 1238 and 1243(1)) is present in all stars hotter than type B 0 for the main sequence, B 1 for the giants and B 2 for the supergiants (Abbottet al. 1982, Dean and Bruhweiler 1985). The line often has a shell component.
The line structure is interpreted as signifying mass loss (Howarth and Prinja 1989). At O4V W(1240) 15, at B 1 Ib, W = 2 Å.
In Be stars the 1240 feature is visible up to B1 - 2V (Marlborough 1982).
Emission lines of NV
For O-type stars and B-type supergiants, see above.
NV is present in at least one pre-degenerate star (Werner et al. 1991). The 1238-1242 lines, when seen in late type stars, indicate the existence of a transition region (see Part Two, section 3.1) characteristic of `coronal stars'. For instance the feature is seen in T Tau stars, together with emission lines of NIV (Appenzeller et al. 1980).
The nitrogen sequence in WR stars
Wolf-Rayet stars with strong nitrogen features are called WN stars. Nitrogen appears in the form of NIII, NIV and NV, with the latter being more prominent in early subtypes (WN 3-WN 6). The most important feature is a strong and broad emission in the region 4600-4640, formed by a blend of NV (4604 and 4620) and NIII (4634, 4640 and 4642) lines. In earlier stars NV predominates, with a subsequent blue shift of the average wavelength of the feature, whereas in later stars one finds mostly N III. The combined feature has an equivalent width between 10 and 100 Åand half-widths of up to 65 Å. Another important emission feature is N IV at 3478-83, with equivalent widths similar to those of the 4600-40 feature. NIV is present in practically all WN stars, with a maximum around WN 6. The N IV emission line 7115 has equivalent widths up to 220 Å(Conti et al. 1990).
Other emission features present in WN stars include HeII 4686 (10 < W < 500) and HeI 5876 (4 < W < 60 Å). For more details see HeI and HeII.
Another strong emission that appears in most WN stars is that of CIV at 5801-12, whose W may reach 1000Å.
In the infrared spectral range the same species appear in emission, with the addition of HI. One thus finds HeII 10124, 8237,6891 and 6683, the first of the four lines being the most intense (W values up to 400 Å) whereas the others are weaker and may be absent. In addition one finds in emission HeI (7065 and 6678), with 6678 more intense than 7065. These HeI lines are, however, not present in all WN stars. The most interesting emission in that of HI, P7 (10049), with W up to 50 Å. Usually H lines are not seen in the blue spectral region, or at best they are seen as weak absorption features, so that the P7 emissions give the most direct information that H is present in WN stars. However, the line is not present in all WN stars (Conti et al. 1990).
In the ultraviolet one also finds the same atomic species in emission, namely NV, NIV, N III and NII, C IV and C III, He II and in addition OV (in the earliest subtypes) and OIV (in the later subtypes) plus SiIV. For a detailed account see Willis et al. (1986).
Behavior in non-normal stars
The lines 4640, 4634 and 4642 of M.2 of NIII appear in emission in supergiants with W < 2 Å. Usually, but not always, the NIII emission is accompanied by HeII 4686, in emission (see He) and by NIV 4058. When N III and He II both appear in emission, the star is called Of When NIII is in emission and HeII in weak emission or in absorption, the star is called O(f) and, if NIII is in emission whereas HeII is in strong absorption, the star is called O((f)). For more details see Jaschek and Jaschek (1987a).
N is about normal in sdO and sdB stars (Heber 1987) but it is strong in He-rich subdwarfs (Husfeld et al. 1989). In HB stars, it is about normal (Heber 1991).
NII and NIII lines are highly enhanced in the spectra of the N variety of the so-called CNO stars (which some authors call OBN), where one finds W values larger by factors of 2-4 than in normal stars of the same spectral type (Schoenberner et al. 1988). In the ultraviolet spectrum of these objects one also sees the NV lines 1239(1) and 1243(5) prominently (Walborn and Panek 1985). In contrast these lines are very weak in the CNO stars with enhanced carbon (Walborn and Panek, 1985).
NI has been investigated in early peculiar stars by Roby and Lambert (1990). They find that NI is slightly weak in Am and Bp stars of the Si subgroup, and much weaker in Hg-Mn and Cr-Eu-Sr stars. In the latter the W values of the NI lines are weakened by factors of the order of four or more with respect to normal stars. NI lines are about normal in lambda Boo stars (Bascheck et al. 1984).
Strong NI lines are present in the spectrum of the A-type He-strong star upsilon Sgr (Greenstein 1943).
N seems to behave in a manner parallel to that of Fe in stars of the Magellanic Clouds (Barbuy et al. 1981, Spite and Spite 1990).
N abundances in late type stars can only be derived through the behavior of the CN molecule - but of course one needs also to analyze other molecules to cope with the C abundance, such as CO, NH and C2. From such analysis it turns out that N is about normal in O-rich giants (M stars and S stars and perhaps slightly underabundant in C stars (Lambert et al. 1986). There exist, however, some S stars that are strong in N (Lambert 1989).
N behaves in a manner parallel to that of Fe in metal-weak stars (Carbon et al. 1987, Wheeler et al. 1989). Curiously, there exist a few metal-weak dwarfs in which N is overabundant (with respect to iron). The percentage of stars with such anomalies is of the order of 4%. Spite and Spite (1986) have shown that, in these stars, Li behaves normally as in other halo dwarfs. Very recently also metal-weak giants have been found in which nitrogen is either abnormally strong or weak with respect to iron (Anthony-Twarog et al. 1992).
Nitrogen is strong in the weak G-band stars (Cottrell and Norris 1978) and in R CrB stars. However, in the latter a large scatter exists among the different stars (Cottrell and Lambert 1982).
In globular cluster giants it has been shown that, at equal Fe/H, there exist variations in CN strength from one cluster to another (Langer et al. 1992). Although most of this is attributed to variations in C, variations in the N strength may also exist.
[NII] lines (like that at 5755) are seen in emission in the principal spectrum phase of novae, together with emission lines of NIII, NIV and NV in the ultraviolet spectral region. During the Orion spectrum phase, absorption lines of NIII (4097 and 4103) and NV (4603 and 4629) are present. At mid-phase of the Orion spectrum there appears a strong 4640 emission, produced by a blend of NIII and NII lines, which persists well into the nebular stage phase. Besides this emission feature one observes also emissions of NII (4995,5680 and 5950), NIII (4097, 4103 and 3484) and NIV(4058). This exceptional enhancement of the N lines is called 'nitrogen flaring'.
At the post-novae phase, [N II] 6584 is seen in emission (Warner 1989, Payne-Gaposchkin 1957). In general N is overabundant by two orders of magnitude in novae (Andreae 1993).
N has two stable isotopes, namely N14 and N15, and six unstable ones. In the solar system the two stable isotopes occur with frequencies of 99.6% (N14) and 0.4% (N15)
Querci and Querci (1970) identified some lines of C12 / N15 in an early C star, which implies a ratio N14 / N15 of several thousands, whereas in the solar system the ratio is about 270.
N14 and N15 are both produced by hydrogen burning and N15 also by explosive hydrogen burning.
Published in "The Behavior of Chemical Elements in Stars", Carlos Jaschek and Mercedes Jaschek, 1995, Cambridge University Press.