do the spectral types DA, DC, etc. mean?
Author: Mike Dworetsky <mmd@star.ucl.ac.uk>

The MK classification system for the vast majority of stars works
remarkably well for one simple reason: most stars in the Galactic disk
have surface chemical compositions that are broadly similar to each
other and the Sun's composition. They are 71 percent hydrogen, 27
percent helium, and 2 percent "metals" (Li--U). Thus, the differences
in spectral line strengths that give rise to the familiar OBAFGKM
sequence are due to their vast range in surface temperature. The MK
system can also classify by absolute stellar brightness: the more
subtle differences in the strengths of certain lines at various
classes, caused by the different surface gravities of main sequence
and supergiant stars, for example, are spoken of as luminosity
criteria, because they depend on the size of the star (big stars
radiate much more energy than small stars, but their atmospheres are
much less dense).

The name "white dwarf" for these stars comes from the observed colors
of the first examples discovered.  They caught the attention of
astronomers because they had large masses comparable to the Sun but
were hot and very faint, hence extremely small and dense.  We now know
that there are a few "white dwarfs" that are actually cool enough to
look red.

The first spectroscopic investigators of white dwarfs tried to fit
them into a descriptive system parallel to the MK classes, using the
letter D plus a suffix OBAFGK or M, with the letter C added for the
cases when the spectra showed no lines (continuous spectra).  The
types were sometimes supplemented by cryptic abbreviations like "wk"
for weak; "s" for sharp-lined, and so on.

When the spectra of white dwarfs were investigated in more detail, it
proved impossible to categorize them neatly for one increasingly
apparent reason: the surface compositions of white dwarfs varied
enormously from star to star.  Astronomers needed a new scheme to
reflect this.  In the revised classification scheme, white dwarf
designations still start with the letter D to indicate dwarf or
"degenerate" stellar structure. A second letter indicates the main
spectral features visible: C for a continuous spectrum with no lines,
A for Balmer lines of hydrogen with nothing else, B for He I (neutral
helium) lines, O for He II with or without He I or H, Z for metal
lines (often, strong Ca II lines are seen), and Q for atomic or
molecular lines of carbon (C is used for continuous spectra; K for
Karbon could be confused with the K stars; so try to think of
Qarbon!).

These basic types can sometimes mix; DAQ stars are known, for example.

A further suffix can be added: P for magnetic stars with polarized
light, H for magnetic stars that do not have polarized light, and V
for variable.  (There is a class of short-period pulsating white
dwarfs, called ZZ Ceti stars.) There may be emission lines (E). And if
an unusual star still defies classification, it goes into type X.

Finally, a number is appended that classifies the star according to
its effective temperature based on formulae which use the observed
colors: the number is 50400/T rounded to the nearest 0.5, i.e., the
value of 50400/temperature, rounded.  If white dwarfs with T much
higher than 50,000 K are ever found, they could have the number 0 or
0.5 appended. The coolest designation is open-ended; there is a star
classified as DC13, for example, which is actually rather red, not
white.

Thus a hot white dwarf with neutral helium lines might be described as
DB2.5; a cooler white dwarf with hydrogen lines, a magnetic field,
polarized light, and a trace of carbon might be DAQP6.

This system can provide good summary descriptions of the vast majority
of white dwarf stars.  However, it is a definite move away from the
original concept of spectral classification, because it requires
photometry and polarimetry as well as visual inspection of a spectrum,
in order to make an assignment.  But most leading experts on the
subject have agreed it was necessary to move in this direction.

Some references:
Sion, E.M., et al. 1983. Astrophys. J., 269, 253--257
Greenstein, J. 1986. Astrophys. J., 304, 334--355
Wesemael, F. et al. 1993. Publ. Astr. Soc. Pacif., 105, 761--778

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