Geology Reference
In-Depth Information
electron in a hydrogen atom, for example, prevails
only as long as the electron's kinetic energy is lower
than the atom's ionization energy (Figure 5.6). Raise
the temperature high enough (to around 10
3
K), and
hydrogen undergoes thermal ionization to form a
plasma of free protons and electrons, although to liber-
ate
inner
electrons from heavier atoms (Chapter 6) the
temperature must rise further (~5 × 10
3
K). At ~10
9
K
even protons and neutrons bound within a nucleus
gain sufficient thermal kinetic energy to surmount the
strong force
and separate from each other. In the cool-
ing, post-Big Bang cosmos these processes operated in
reverse, as particles combined for the first time into
nuclei and ultimately atoms.
process was first proposed - though based on earlier
insights by Fred Hoyle - remains the foundation of our
understanding of element formation today.
We can visualize stellar nucleosynthesis as a long
series of consecutive steps, like an industrial assembly
line. Not every stellar factory, however, possesses the
full assembly line. Fusion reactions in stars take place
in a series of stages, hand in hand with the thermal
evolution of the star, and how far the process may go
depends, as we shall see, on the mass of the star.
Hydrogen is consumed to form helium early in
the development of a star (see Table 11.1). When
hydrogen in the centre of a star (where these reactions
occur most quickly owing to the high density and
temperature) is almost used up, the star raises its core
temperature (by gravitational contraction) to a level
sufficient to allow helium nuclei to combine to form
carbon and oxygen. Similarly helium must more or
less run out in the core of the star before further
contraction and heating can occur, allowing carbon
and oxygen to be transformed to heavier elements
leading to silicon.
The maximum temperature a star can achieve
during normal evolution is related to its mass. A star
of the Sun's mass (
M
⊙
) is capable only of stages 1 and
2. A star probably needs to have a mass exceeding
30
M
⊙
before all fusion reactions leading to iron
become possible (Tayler, 1975). Even such massive
stars only generate heavy elements like iron at their
very centres. Many stars fall below this mass range,
and therefore contribute only to the abundance of the
lighter elements. The general fall-off in abundance
towards heavier nuclides (item (b) on p. 210) reflects
the relatively small number of stars capable of
generating the heaviest elements.
Fusion reactions can generate most but not all of the
stable nuclides between hydrogen and iron. As
8
Be is
a very unstable nucleus, the main fusion reactions
evidently proceed directly from
4
He to
12
C, largely
bypassing the elements Li, Be and B. The small amounts
of Be and B shown in Figures 11.2 and 11.3 - item (c) on
p. 210 - actually seem to have been produced by the
breakdown of heavier nuclei (
12
C,
16
O) under cosmic-
ray bombardment, a process called
spallation
.
The balance of nuclear forces gives nuclei in the iron
mass range the greatest stability (Box 11.2). Massive
stars having sufficiently high temperatures can pro-
duce these nuclides relatively efficiently, hence the 'iron
Stars
These light nuclides formed in the Big Bang provided
the feedstock for the manufacture of all the heavier
chemical elements that make up the Earth and the
other baryonic matter in the cosmos. This manufactur-
ing process has been going on throughout the life of
the Universe. Evidently it has not been very 'efficient',
because
1
H and
4
He still make up 98% of the mass of
the observable baryonic universe. How has this small
inventory of heavier elements been formed?
The principal process for generating the elements up
to iron is nuclear fusion (Box 11.2). Nuclear fusion can
only occur if two conditions are both satisfied:
(i) a high density of matter to raise the probability of
nuclear collisions (which in interstellar space is
vanishingly small); and
(ii) a high temperature (at least 10
7
K) to ensure that
positively charged nuclei will collide with sufficient
kinetic energy to overcome their mutual electro-
static repulsion; nuclei need to approach closer than
10
-14
m before the
strong force
begins to operate to
bind them together into a heavier nucleus.
The interior of a star furnishes both of these require-
ments, and the abundances of heavier elements that
we see today are regarded as the cumulative prod-
uct of nucleosynthesis that has taken place inside
many generations of stars. Despite much subsequent
research, the visionary 1957 paper by Burbidge,
Burbidge, Fowler and Hoyle ('B
2
FH') in which this
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