poses Bd ele ote Bee Da
wo ee at a
outermost layers of the model. In this
case however. one may be certain that
the equation of transfer that has been
adopted to represent line formation is
star different from the one seen in a
is not appropriate to calculate degree
of ionization and level of excitation
by the equilibrium relations known as
Saha’s and Boltzmann’s laws. In solv-
the brightness at 4500 A, is due to the
no longer tenable: furthermore, that it
ing the problems of the monochromatic
transfer of radiation through the outer
parts of the stellar atmosphere one
should take into account what are
called “non-lt.e. effects’—that is, departures from the equations represent-
ing local thermodynamic equilibrium.
In order to obtain a first prediction
brightness (Fig.
the
spectrum has been predicted by use of
the existing thermal-equilibrium theory
of line formation. When observed line
profiles are compared with predicted
line profiles, attention must be paid to
isolation of discrepancies due to (i) de-
partures from the adopted simple theory of line formation: (ii) departures of
the state of motion, of the gas in the
plane parallel layers, from the thermal
motions appropriate to the local temperature, which are the only motions
bly deeper layers than does the spectrum
Expected Ultraviolet Absorption Lines
tion of the depth of an absorption-
The ultraviolet spectrum of O and
B stars divides into two regions. In
the section 911.6 to 1900 A there
at 4500 A. Important in determina-
line profile at any spectral region is
the factor of whether the selected
frequency range lies on the steep side
the
stronger
mosphere, and where it occurs, is the
“depth of formation” of the spectrum
at frequency ». Whether this frequency lies in the continuous spectrum,
Or in
shown that Fp, the emergent flux in
frequency y, is approximately equal
to By, the value of the Planck function at this frequency, at the layer
in the model in which the monochro-
matic optical depth t) is about equal
to 0.4. In a very opaque spectral region the position at which ftequals
0.4 will occur in the outermost layers
of the star; in a transparent spectral
region fy will equal 0.4 at much greater depths. In the first case, the local
temperature will be low; in the second
it may be higher by 10,000° or 15.000°K. Clearly, in an opaque spectral
region such as the centre of a strong
line, one is observing a part of the
8 DECEMBER 1967
lines
are
not
deep;
at
only important spectra not having reso-
nance lines in this region are those of
He I and H, O II and
Il, and Mg II. Morton
dicted the profiles of
stronger lines formed in
cially for resonance lines, where line-
formation sketched above. The resuits
that the presence of strong lines in the
in the neighborhood of 1200 A appear
in Fig. 3; the great width and depth
ultraviolet spectral region causes greater reduction of the expected emergent
flux, from the flux for the case of no
a
We
ee
es
ee
ee
Fy
~
of the resonance lines are evident. The
central parts of these lines are formed
ee
e
ee
e
ee
ee es
UNBLANKETED
rh.
'
ee
ee
ee
"
———=
MODEL ----
.
TYPE B15
rs
it
ee
PREDICTED SPECTRUM
LINE-BLANKETED MODEL
4o0b
IH, Ne I and
(4) has presome of the
an unblanket-
ed model atmosphere of type about
BI.5, using the simple theory of line
formation by scattering processes must
be important), but it does demonstrate
‘
I
J
-_
hI
3.0}
~
’
{
a line, is immaterial; what is
significant is that numerical work has
second, and third ions of the elements
He to Ca; many of these lines are
predicted to be exceedingly strong. The
cially at very short wavelengths.
The foregoing qualitative explanation does not tell the full story (espe-
An important concept for visualiza-
tion of what occurs in the stellar at-
lines from low-lying levels of the first,
tail. In the latter case (at wavelengths
(iii) the fact that the stellar atmo-
metric form.
occur Many resonance lines and strong
of the Planck function for the relevant
temperatures or on the long-wavelength
taken into account in the theory: and
sphere may indeed be so extensive that
it cannot be adequately represented by
a plane parallel layer, but must be
represented as a sphere or other geo-
shown in Fig. 2.
trum at 1300 A comes from considera-
wavelengths shorter than 3000 A, the
strong lines become very deep, espe-
and
factors causes
fact that in these calculations the spec-
erations have
ignored
play of these various
the two different spectral distributions
1), in comparison to
greater than 3600 A in B stars) even
been
tween 4000 and 5000 A. The inter-
lines. A large part of the ultraviolet
of the ultraviolet spectrum from OB
stars, these important physical consid-
lines, than would ever occur in the
normally observed spectral region be-
transparent spectral region, such as in
the continuous spectrum or in weak
t
— 4
—
i
[
t
2.0+!
i
i
~
+
-
'
olla
|
BALMER
JUMP
ee
-
1000
|
of
poy
tt
2000
lg
tl ng
3000
4000
FG
5000
yl
tl
§ 000
Wavelength (A)
Fig. 2. The difference in spectral distribution between a line-blanketed model calculated
by C. Guillaume and an unblanketed model calculated by one of us (A.B.U.). The ynjt
of flux is 10“ erg per square centimenter per second per unit frequency interyayy:,
{275