Discovering the Solar System
Author: Barrie W. Jones
ISBN: 9780470018309
Publisher: Wiley, John & Sons, Incorporated
Sample Page
Chapter One
The Sun and its Family
Imagine that you have travelled far into the depths of space. From your
distant vantage point the Sun has become just another star amongst the
multitude, and the Earth, the other planets, and the host of smaller bodies
that orbit the Sun are not visible at all to the unaided eye. The Sun is by
far the largest and most massive body in the Solar System, and is the only
one hot enough to be obviously luminous. This chapter starts with a
description of the Sun. We shall then visit the other bodies in the Solar
System, but only briefly, the purpose here being to establish their main
characteristics - each of these bodies will be explored in much more detail
in subsequent chapters. Chapter 1 then continues with an exploration of the
orbits of the various bodies. Each of them also rotates around an axis
through its centre, and we shall look at this too. The chapter concludes
with aspects of our view of the Solar System as we see it from the Earth.
1.1 The Sun
This is only a very brief account of the Sun, and it is biased towards
topics of importance for the Solar System as a whole. Fuller accounts of the
Sun are in books listed in Further Reading.
1.1.1 The Solar Photosphere
The bright surface of the Sun is called the photosphere (Plate 1).
Its radius is 6 96 x [10.sup.5] km, about 100 times the radius of the Earth.
It is rather like the 'surface' of a bank of cloud, in that the light
reaching us from the photosphere comes from a range of depths, though the
range covers only about one-thousandth of the solar radius, and so we are
not seeing very deep into the Sun. It is important to realise that whereas a
bank of cloud scatters light from another source, the photosphere is
emitting light. It is also emitting electromagnetic radiation at other
wavelengths, as the solar spectrum in Figure 1.1 demonstrates. The total
power radiated is the area under the solar spectrum, and is 3 85 x
[10.sup.26] watts (W). This is the solar luminosity. The photosphere, for
all its brilliance, is a tenuous gas, with a density of order [10.sup.-3] kg
[m.sup.-3], about 1000 times less than that of the air at the Earth's
surface.
The spectrum in Figure 1.1 enables us to estimate the mean photospheric
temperature. This is done by comparing the spectrum with that of an ideal
thermal source, sometimes called a black body. The exact nature of such
a source need not concern us. The important point is that its spectrum is
uniquely determined by its temperature. Turning this around, if we can fit
an ideal thermal source spectrum reasonably well to the spectrum of any
other body, then we can estimate the other body's temperature. Figure 1.1
shows a good match between the solar spectrum and the spectrum of an ideal
thermal source at a temperature of 5770 K. Also shown is the poor match with
an ideal thermal source at 4000 K, where the peak of the spectrum is at
longer wavelengths. Also, the power emitted by this source is a lot less.
The power shown corresponds to the assumption that the 4000 K source has the
same area as the source at 5770 K, and thus brings out the point that the
temperature of an ideal thermal source determines not only the wavelength
range of the emission, but the power too. Note that 5770 K is a
representative temperature of the Sun's photosphere; the local
temperature varies from place to place.
At a finer wavelength resolution than in Figure 1.1 the solar spectrum
displays numerous narrow dips, called spectral absorption lines. These are
the result of the absorption of upwelling solar radiation by various atoms
and ions, mainly in the photosphere, and therefore the lines provide
information about chemical composition. Further information about the Sun's
composition is provided by small rocky bodies that continually fall to
Earth. They are typically 1-100 cm across, and constitute the meteorites
(Section 3.3). At 5770 K significant fractions of the atoms of some elements
are ionised, and so it is best to define the composition at the photosphere
in terms of atomic nuclei, rather than neutral atoms. In the photosphere,
hydrogen and helium dominate, with hydrogen the most abundant - all the
other chemical elements account for only about 0.2% of the nuclei. Outside
the Sun's fusion core (Section 1.1.3) about 91% of the nuclei are hydrogen
and about 9% are helium.
Plate 1 shows that the most obvious feature of the photosphere is dark
spots. These are called (unsurprisingly) sunspots. They range in size
from less than 300 km across to around 100 000 km, and their lifetimes range
from less than an hour to 6 months or so. They have central temperatures of
typically 4200 K, which is why they look darker than the surrounding
photosphere. Sunspots are shallow depressions in the photosphere, where
strong magnetic fields suppress the convection of heat from the solar
interior, hence the lower sunspot temperatures. Their number varies,
defining a sunspot cycle. The time between successive maxima ranges from
about 8 years to about 15 years with a mean value of 11.1 years. From one
cycle to the next the magnetic field of the Sun reverses. Therefore, the
magnetic cycle is about 22 years.
Sunspots provide a ready means of studying the Sun's rotation, and reveal
that the rotation period at the equator is 25.4 days, increasing with
latitude to about 36 days at the poles. This differential rotation is common
in fluid bodies in the Solar System.
1.1.2 The Solar Atmosphere
Above the photosphere there is a thin gas that can be regarded as the solar
atmosphere. Because of its very low density, at most wavelengths it emits
far less power than the underlying photosphere, and so the atmosphere is not
normally visible. During total solar eclipses, the Moon just obscures the
photosphere, and the weaker light from the atmosphere then becomes visible.
In Plate 2 the atmosphere just above the photosphere is not visible, whereas
in Plate 3 the short exposure time has emphasised the inner atmosphere. The
atmosphere can be studied at other times, either by means of an optical
device called a coronagraph that attenuates the radiation from the
photosphere, or by making observations at wavelengths where the atmosphere
is brighter than the photosphere.
Figure 1.2 shows how the temperature and density in the solar atmosphere
vary with altitude above the base of the photosphere. A division of the
atmosphere into two main layers is apparent, the chromosphere and the
corona, separated by a thin transition region.
The chromosphere
The chromosphere lies immediately above the photosphere. It has much the
same composition as the photosphere, so hydrogen dominates. The density
declines rapidly with altitude, but the temperature rises. The red
colour that gives the chromosphere its name ('coloured sphere') is a result
of the emission by hydrogen atoms of light at 656.3 nm. This wavelength is
called H[alpha] 'aitch-alpha').
The data in Figure 1.2 are for 'quiet' parts of the chromosphere. Its
properties are different where magnetic forces hold aloft filamentary clouds
of cool gas, extending into the lower corona. The filaments are the red
prominences above the limb of the photosphere in Plate 3. Prominences are
transitory phenomena, lasting for periods from minutes to a couple of
months. The chromosphere is also greatly disturbed in regions where a flare
occurs. This is a rapid brightening of a small area of the Sun's upper
chromosphere or lower corona, usually in regions of the Sun where there are
sunspots. The increase in brightness occurs in a few minutes, followed by a
decrease taking up to an hour, and the energy release is spread over a very
wide range of wavelengths. Flares, like certain prominences, are associated
with bursts of ionised gas that escape from the Sun. Magnetic fields are an
essential part of the flare process, and it seems probable that the
electromagnetic radiation is from electrons that are accelerated close to
the speed of light by changes in the magnetic field configuration. As with
so many solar phenomena, the details are unclear.
The corona
Above the chromosphere the density continues to fall steeply across a thin
transition region that separates the chromosphere from the corona (Figure
1.2).
* What distinctive feature of the transition region is apparent in Figure
1.2?
A distinctive feature is the enormous temperature gradient. This leads into
the corona, where the gradient is not so steep. The corona extends for
several solar radii (Plate 2), and within it the density continues to fall
with altitude, but the temperature continues to rise, reaching 3-4 x
[10.sup.6] K, sometimes higher. Conduction, convection, and radiation from
the photosphere cannot explain such temperatures - these mechanisms would
not transfer net energy from a body at lower temperature (the
photosphere) to a body at higher temperature (the corona). The main
heating mechanism seems to be magnetic - magnetic fields become reconfigured
throughout the corona, and induce local electric currents that then heat the
corona. Waves involving magnetic fields (magnetohydrodynamic waves) also
play a role in certain regions.
The corona is highly variable. At times of maximum sunspot number it is
irregular, with long streamers in no preferred directions. At times of
sunspot minimum, the visible boundary is more symmetrical, with a
concentration of streamers extending from the Sun's equator, and short,
narrow streamers from the poles. Coronal 'architecture' owes much to solar
magnetic field lines. The white colour of the corona is photospheric light
scattered by its constituents. Out to two or three solar radii the
scattering is mainly from free electrons, ionisation being nearly total at
the high temperatures of the corona. Further out, the scattering is
dominated by the trace of fine dust in the interplanetary medium.
The solar wind
The solar atmosphere does not really stop at the corona, but extends into
interplanetary space in a flow of gas called the solar wind, which
deprives the Sun of about one part in 2 5 x [10.sup.-14] of its mass per
year. Because of the highly ionised state of the corona, and its
predominantly hydrogen composition, the wind consists largely of protons and
electrons. The temperature of the corona is so high that if the Sun's
gravity were the only force it would not be able to contain the corona, and
the wind would blow steadily and uniformly in all directions. But the strong
magnetic fields in the corona act on the moving charged particles in a
manner that reduces the escape rate. Escape is preferential in directions
where the confining effect is least strong, and an important type of
location of this sort is called a coronal hole. This is a region of
exceptionally low density and temperature, where the solar magnetic field
lines reach huge distances into interplanetary space. Charged particles
travel in helical paths around magnetic field lines, so the outward-directed
lines facilitate escape. The escaping particles constitute the fast wind.
Elsewhere, where the field lines are confined near the Sun, there is an
additional outward flow, though at lower speeds, called the slow wind.
Solar wind particles (somehow) gain speed as they travel outwards, and at
the Earth the speeds range from 200 to 900 km [s.sup.-1]. The density is
extremely low - typically about 4 protons and 4 electrons per [cm.sup.3],
though with large variations. Particularly large enhancements result from
what are called coronal mass ejections, often associated with flares and
prominences, and perhaps resulting from the opening of magnetic field lines.
If the Earth is in the way of a concentrated jet of solar wind, then various
effects are produced, such as the aurorae (the northern and southern lights
- Plate 26). The solar wind is the main source of the extremely tenuous gas
that pervades interplanetary space.
Solar activity
Solar activity is the collective term for those solar phenomena that
vary with a periodicity of about 11 years.
* What two aspects of solar activity were outlined earlier?
You have already met the sunspot cycle, and it was mentioned that the form
of the corona is correlated with it. Prominences (filaments) and flares are
further aspects of solar activity, both phenomena being more common at
sunspot maximum. The solar luminosity also varies with the sunspot cycle,
and on average is about 0.15% higher at sunspot maximum than at
sunspot minimum. This might seem curious, with sunspots being cooler and
therefore less luminous than the rest of the photosphere. However, when
there are more sunspots, a greater area of the photosphere is covered in
bright luminous patches called faculae.
All the various forms of solar activity are related to solar magnetic fields
that ultimately originate deep in the Sun. The origin of these fields will
be considered briefly in the following description of the solar interior.
1.1.3 The Solar Interior
To investigate the solar interior, we would really like to burrow through to
the centre of the Sun, observing and measuring things as we go. Alas! This
approach is entirely impractical. Therefore, the approach adopted, in its
broad features, is the same as that used for all inaccessible interiors. A
model is constructed and varied until it matches the major properties that
we either can observe, or can obtain fairly directly and reliably
from observations. Usually, a range of models can be made to fit, so
a model is rarely unique. Many features are, however, common to all models,
and such features are believed to be correct. This modelling process will be
described in detail in Chapter 4, in relation to planetary interiors. Here,
we shall present the outcome of the process as applied to the Sun.
A model of the solar interior
Figure 1.3 shows a typical model of the Sun as it is thought to be today.
Hydrogen and helium predominate throughout, as observed in the photosphere.
Note the enormous increase of pressure with depth, to [10.sup.16] pascals
(Pa) at the Sun's centre - about [10.sup.11] times atmospheric pressure at
sea level on the Earth! The central density is less extreme, 'only' about 14
times that of solid lead as it occurs on the Earth, though the temperatures
are so high that the solar interior is everywhere fluid - there are no
solids. Another consequence of the high temperatures is that at all but the
shallowest depths the atoms are kept fully ionised by the energetic atomic
collisions that occur. A highly ionised medium is called a plasma.
The central temperatures in the Sun are about 1 4 x [10.sup.7] K,
sufficiently high that nuclear reactions can sustain these temperatures and
the solar luminosity, and can have done so for the 4600 million years (Ma)
since the Sun formed (an age based on various data to be outlined in Chapter
3, notably data from radiometrically dated meteorites). This copious source
of internal energy also sustains the pressure gradient that prevents the Sun
from contracting.
Though nuclear reactions sustain the central temperatures today, there must
have been some other means by which such temperatures were initially
attained in order that the nuclear reactions were triggered. This must have
been through the gravitational energy released when the Sun contracted from
some more dispersed state. With energy being radiated to space only from its
outer regions, it would have become hotter in the centre than at the
surface. Nuclear reaction rates rise so rapidly with increasing temperature
that when the central regions of the young Sun became hot enough for nuclear
reaction rates to be significant, there was a fairly sharp boundary between
a central core where reaction rates were high, and the rest of the Sun where
reactions rates were negligible. This has remained the case ever since. At
present the central core extends to about 0.3 of the solar radius (Figure
1.3). This is a fraction [(0.3).sup.3] of the Sun's volume, which is only
2.7%. However, the density increases so rapidly with depth that a far
greater fraction of the Sun's mass is contained within its central core.
(Continues...)