The Sun
Our Star
© Sierra College Astronomy Department
The Sun - Our Star
Solar Data


As viewed from the Earth, the Sun
has an average angular diameter of
31’ 59” and is at an average distance
from the Earth of 1.50 X 108 km.
The diameter of the Sun can therefore
be calculated as 1.39 X 106 km, or
about 110 times Earth’s diameter and
about 10 times Jupiter’s.
© Sierra College Astronomy Department
2
The Sun - Our Star
Solar Data




The mass of the Sun is 1.99 X 1030 kg,
more than 300,000 times Earth’s mass.
The Sun’s average density is 1.41 g/cm3,
about the same as the density of Jupiter.
The escape velocity of the Sun is 620 km/s.
The Sun rotates in 25.4 days at its equator
and 36 days near its poles.
© Sierra College Astronomy Department
3
The Sun - Our Star
Solar Data





The Sun emits energy is all portions of the
electromagnetic spectrum with its peak in the
visible portion.
Luminosity is the total power output of the Sun,
which is equal to 3.8 X 1026 watts.
Brightness: Solar energy strikes the Earth at the
rate of 1,380 watts/m2 (Solar Constant).
We can examine the surface of the sun by looking
at its spectra (e.g., the Fraunhofer lines).
The Sun is a G2V star with an absolute
magnitude of 4.8.
© Sierra College Astronomy Department
4
The Sun - Our Star
Basic Solar Physics


To estimate the internal parameters of
the Sun, astronomers must solve a set
of equations that describe how pressure,
temperature, mass, and luminosity
change.
We need to know the chemical
composition, rate of energy production
and surface conditions to solve these
equations.
© Sierra College Astronomy Department
5
The Sun - Our Star
Solar Energy Source
Must produce 3.8 x 1026 Watts
 Combustion
 Hydrogen + Oxygen produces energy but would last only
1000 years

Meteorite Impacts
 The infall of material can be converted into heat and light
 About 1/10 Earth masses of debris would have to fall into the
Sun each year

Gravitational contraction
 The infall of material produces energy
 Could sustain sun for about 500 million years
 Still short of geologic age of Earth (~ 4.5 billion yr)
© Sierra College Astronomy Department
6
The Sun - Our Star
Solar Energy Source

Mass/Energy Conversion
 Einstein proposed in 1905 that mass and
energy are interconvertible - derives famous
E = mc2 equation.
 Nuclear Fission
Large atoms break into two and release energy
 Problem: Sun does not contain enough of the big atoms

 Nuclear Fusion
Light atoms combine to produce larger atom and release
energy
 Sun full of small atoms
 Turns out to be the right answer

© Sierra College Astronomy Department
7
The Sun - Our Star
Solar Energy Source




In nuclear fusion, two nuclei (consisting of protons
and neutrons) combine to form a larger nucleus,
releasing energy in the process.
Large temperatures (15.6 million K) and densities
(150,000 kg/m3) are needed to overcome
electromagnetic repulsion between the protons.
This only occurs in the Sun’s core (inner 10% of
its radius).
In the Sun, 4 hydrogen nuclei are fused to form 1
helium nucleus and energy (proton-proton chain).
© Sierra College Astronomy Department
8
The Sun - Our Star
Solar Energy Source




To produce the Sun’s energy output nearly
4 million tons of matter must be converted
into energy each second.
This in turn requires 610 billion kg of
hydrogen be transformed into 606 billion kg
of helium.
Solar lifetime: ~ 10 billion years
But how do we confirm this from
observations?
© Sierra College Astronomy Department
9
The Sun - Our Star
Solar Neutrinos




Neutrinos are highly non-interacting particles
which can be used as a direct probe into the
Sun (if our nuclear fusion idea is right)
However, they are incredibly hard to detect
Huge underground detectors are used to look
for neutrinos from the Sun
Until recently there seemed to be too few
coming from the Sun
 Solution: neutrinos have a little mass and come in
different “flavors” and change identity on the way to
Earth!
© Sierra College Astronomy Department
10
The Sun - Our Star
Basic Interior Structure
Hydrostatic Equilibrium
 In a star or a planet, hydrostatic equilibrium
is the balance between pressure caused by
the weight of material above and the
upward pressure exerted by material
below.
 Pressure at the Sun’s center is calculated
to be 1.3 X 109 times that on the surface of
the Earth.
© Sierra College Astronomy Department
11
The Sun - Our Star
Basic Interior Structure


Helioseismology studies the Sun’s interior by
observing pulsations in the photosphere.
Many patterns or modes of oscillation can be
seen and resemble those from a string on a
violin or seismic waves from earthquakes.
 Typical speeds of up/down motions are 100 to 300
m/s with periods of 5 minutes.
 Short wavelength modes of a few 1000 km
resonate only in outermost parts of Sun. Longer
wavelengths penetrate deeper.
© Sierra College Astronomy Department
12
The Sun - Our Star
Basic Interior Structure

Helioseismology has been used to detect:
 The bottom of the convection zone
 The Sun has more helium than previously thought
 The interior has surprisingly about the same
rotation period as the surface
 An extensive gas stream system beneath the
surface and as much as 30% of the way to the
Sun’s center
 Sunspots on the far side of the Sun
© Sierra College Astronomy Department
13
The Sun - Our Star
Energy Transport


There are three possible mechanisms
for energy transport from one location
to another: conduction, radiation, and
convection.
Conduction is the transfer of energy
within a substance by collisions
between atoms and/or molecules. This
is not a significant factor in transporting
energy within the Sun.
© Sierra College Astronomy Department
14
The Sun - Our Star
Energy Transport
Radiation
 The energy produced by the Sun is carried out by
photons emitted at one spot and absorbed at another
(radiative diffusion).
 The speed at which photons get out from the Sun
depends heavily on its opacity – the ability of a
substance to stop photons (the opposite of
transparency).
 In the Sun’s core, the photons typically travel 10-6 m
before being reabsorbed.
 About 1025 absorptions and reemissions are needed
before the energy reaches the Sun’s surface.
© Sierra College Astronomy Department
15
The Sun - Our Star
Energy Transport
Convection
 Radiative diffusion carries solar energy out to 70% of
the Sun’s radius.
 At this distance, the temperature has dropped to 1.5
million K and hydrogen atoms start to from, increasing
the opacity.
 As a result, the rate of temperature decline becomes
steep and convection takes over the energy transport
in the outer 30% of the Sun’s radius
 Typically, the diffusion of energy from core to surface
takes on average 170,000 years.
© Sierra College Astronomy Department
16
The Sun - Our Star
The Photosphere

Photosphere is the visible “surface” of the Sun.
 It is the part of the solar atmosphere from which light
(most energy is in optical) is emitted into space.
 The photosphere is a very thin layer - 200 km thick.


The photosphere varies in temperature from about
8,000 K at its deepest to 4,000 K near its outer edge.
 Overall, the light received from the photosphere is
representative of an object about 5,800 Kelvin.
Pressure of the outer photosphere is 0.01 the pressure of the
Earth’s surface. The calculated density of particles is 0.001 of
the density of air at sea level.
© Sierra College Astronomy Department
17
The Sun - Our Star
The Photosphere

Limb is the apparent edge of the Sun
(or any object) seen in the sky.
 The edge of the Sun appears dimmer
than the center. This is known as limb
darkening.
 Observing the limb of the Sun one sees
to a lesser depth because the line of
sight is at a grazing angle.
© Sierra College Astronomy Department
18
The Sun - Our Star
The Photosphere

Granulation is the division of the Sun’s surface
into small convection cells.
 Granules are areas where hot material (light areas)
is rising from below and then descending (dark
surroundings).
 Each granule is typically 1000 km across,
separated by regions 100 K cooler, and last about
15 minutes.

Supergranulation is convective pattern on a much
larger scale, typically 30,000 km across, lasting
about 1 day, and extending to the bottom of the
convection zone.
© Sierra College Astronomy Department
19
The Sun - Our Star
The Photosphere


Mass of the photosphere is 71% hydrogen,
with helium comprising most of the
remainder and a few percent consisting of
several elements found on Earth.
From our knowledge of nuclear fusion, we
know the Sun’s core must hold more
helium. Calculations show that the
hydrogen makes up only 34% of the
center.
© Sierra College Astronomy Department
20
The Sun - Our Star
The Lower Atmosphere



The chromosphere is the region of the
solar atmosphere some 2,000 to 3,000 km
thick that lies between the photosphere
and the corona.
It is not usually observable from Earth
except during a total solar eclipse.
When seen, it has a reddish color, which is
caused by the hydrogen Balmer a emission
line.
© Sierra College Astronomy Department
21
The Sun - Our Star
The Lower Atmosphere

The chromosphere is mostly empty with
the exception of spicules, which:
 Are narrow jets of gas originating at the base of
the chromosphere, shooting upward with
speeds of about 25 km/sec, and lasting about 5
minutes
 form a network above the edges of the
photospheric supergranules
© Sierra College Astronomy Department
22
The Sun - Our Star
The Lower Atmosphere

The corona is the outermost portion of the Sun’s
atmosphere that can only be seen during a total solar
eclipse.
 Much of the light of the corona originates in the photosphere
– light which is then scattered by electrons in the corona.
 Some light is also the result of emission from ionized
elements.


The temperature change from the chromosphere to the
corona is quite dramatic rising from about 4000 K to 1
million K in a transition zone only a few thousand km thick.
Heating of the chromosphere and especially the corona is
not understood and remains one of the unanswered
questions about the Sun (although the heating process
most likely involves energy transport from the convection
zone via the Sun’s time-varying magnetic fields).
© Sierra College Astronomy Department
23
The Sun - Our Star
The Lower Atmosphere
Coronal Structure
 Best seen in X-ray and ultraviolet images
 Bright areas, also referred to as active
regions, are closed magnetic field regions
with hot gas often seen along coronal
loops
 Darker and cooler areas are coronal
holes, which are open magnetic field
regions emitting high-speed gas into
interplanetary space
© Sierra College Astronomy Department
24
The Sun - Our Star
The Lower Atmosphere
Prominences




Dense, relatively cool regions reaching upwards of
50,000 km or more into corona and lasting 2-3 months
Appear as dark filaments against brighter disk of Sun.
Quiescent prominences located away active regions
live longer than active prominences above active
regions.
When a prominence reaches about 50,000 km it
erupts sending gas into interplanetary space with
speeds in excess of 1000 km/s. These eruptions are
called coronal mass ejections.
© Sierra College Astronomy Department
25
The Sun - Our Star
The Lower Atmosphere
Flares





Solar flares, much more explosive and energetic than
prominences, are erupt releases of magnetic energy.
Flares take place in active regions where a prominence is
supported against gravity by magnetic field lines and then the
magnetic field structure changes abruptly.
Coronal gas may heat to 40 million K and X-rays and
ultraviolet light are emitted. The Sun’s brightness may
increase by 1% during an unusually bright flare.
Flares blast out large numbers of very energetic charged
particles that will take about 3 days to reach Earth.
Large solar flares cause spectacular auroras and can affect
earthly radio transmissions if the ionosphere is disrupted by
high-energy particles.
© Sierra College Astronomy Department
26
The Sun - Our Star
The Solar Wind and Heliosphere

The Solar Wind
 The flow of coronal gas into interplanetary space is called
the solar wind.
 Coronal gas takes about 4 days to reach the Earth.
 Near the Earth the solar wind travels at 450 km/s; density
is from 2–10 particles/cm3.
 Speeds in excess of 1000 km/sec are possible during
coronal mass ejections.

The Heliosphere
 The boundary between the region dominated by the
solar magnetic field and interstellar space is called the
heliopause.
 The region interior to the heliopause is the heliosphere.
© Sierra College Astronomy Department
27
The Sun - Our Star
Sunspots and the Solar Activity Cycle





Dark spots on the Sun were first reported by the
Chinese in the 5th century B.C.
Galileo and Thomas Harriott were the first
Europeans to report these sunspots in the early
17th century.
In 1851, Schwabe discovered the sunspot cycle,
which lasts about 11 years.
This periodic cycle, however, is not always
present as the Maunder Minimum appears to
show.
Additionally, individual sunspots are temporary
phenomena lasting from a few hours to a few
months.
© Sierra College Astronomy Department
28
The Sun - Our Star
Sunspots and the Solar Activity Cycle

Magnetic fields can be measured using the
Zeeman effect, which is the splitting of
spectral lines by a strong magnetic field.
Spectra
No magnetic Field

Magnetic Field
Sunspots involve the Sun’s magnetic field The magnetic field in a sunspot is about 1,000
times that of the surrounding photosphere.
© Sierra College Astronomy Department
29
The Sun - Our Star
Why are sunspots dark?



Sunspots are dark because they are cooler than there
surrounding photosphere.
They are cooler because the magnetic field
associated with the sunspot disrupts the convection in
the granule
.
 Heat does not flow as efficiently to the photosphere
 Spot becomes cooler than normal (it is seen as dark, though
it would be red if seen in isolation)

Sunspots, which often appear in pairs aligned in
an east-west direction, have opposite magnetic
polarities - one being north and the other south.
© Sierra College Astronomy Department
30
The Sun - Our Star
Sunspots and the Solar Activity Cycle

Details of the Solar Sunspot Cycle
 At a sunspot maximum, most spots occur
about 30° north or south of the equator.
 As the sunspot cycle progresses, the spots are
seen closer and closer to the Sun’s equator.
 When the spots reach the equator, the cycle is
at a sunspot minimum and begins again.
 Location and relative number of sunspots can
be plotted on a butterfly diagram.
© Sierra College Astronomy Department
31
The Sun - Our Star
Sunspots and the Solar Activity Cycle
A Model for the Sunspot Cycle
 The current model (The Babcock model) has patterns
of magnetic field lines within the Sun’s interior forming
tubes that gradually become twisted from the Sun’s
differential rotation.
 Forced to the surface, these magnetic tubes become
visible as sunspots.
 Breaking the surface weakens the field lines and the
sunspots die out.
 As new lines form deep within the Sun, the magnetic
field direction in the emerging tubes will reverse over
time, causing the magnetic field of the Sun to reverse
in direction in a 22-year period. This is called the
magnetic cycle of the sun.
© Sierra College Astronomy Department
32