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We keep an eye out for solar activity, which can cause Juno and other spacecraft to
experience upsets, e.g., single event upsets or SEUs, which can flip the bits in computer
chips, causing the spacecraft to enter a safe condition as a precaution. Also, some of our
instruments are capable of monitoring solar wind activity in the Sun's "heliosphere"
(since they're designed to measure activity in the magnetosphere of Jupiter), so they can
tell when we've experienced effects from something like a coronal mass ejection (CME)
from the Sun and correlate it with unusual behavior in our onboard computers. We've
tried to correlate the single-bit errors and multi-bit errors we've experienced in flight
with space weather, but unfortunately those events usually seem to occur more
randomly. An alternative explanation for some of our computer glitches is galactic
cosmic rays (GCRs), which are high-energy particles that come from outside our solar
system, but those can come from any direction and aren't really "local weather".
Stuart Stephens
Mission Planner, Juno
Chapter 4: The Sun’s influence on the Earth and other planets
One of the fundamental problems with outer space law, and, yes, there is such a thing, is
that there is no generally accepted definition of “space” (or “outer space”).
The most common definition is “where the atmosphere ends, and the hard vacuum of
space takes over”. The difficulty with this is that there is not a line at which the
atmosphere ends. In fact, there is another definition suggesting space begins where
Earth’s weather ends, but there is such a thing as “space weather”.
The term “space weather” may seem odd; after all, even interplanetary space is a nearvacuum and there is very little water to precipitate! The important idea in the previous
sentence is “near-vacuum”, implying there are at least a few particles out there. The Sun
plays a major role in generating and directing these particles, and their interaction with
planetary atmospheres drives the phenomenon of space weather.
How those particles leave the Sun, get aimed toward the planets and interact with
planetary atmospheres and magnetic fields is the story of this chapter. We begin with the
Sun’s atmosphere as the source of all these partcles.
The Sun’s atmosphere is filled with ionized particles
The structure of the Sun’s interior was described in the previous chapter: core, radiative
zone, convective zone. The surface of the Sun is the outer edge of the convective zone,
and is called the photosphere. This is what we see when we observe the Sun (safely,
through welding glasses or a Sun filter). This is the coolest part of the Sun with a
temperature of 6000 K, which is still hot enough to glow.
The convection in the Sun’s convective zone organizes itself into convection cells, each
cell approximately the size of Texas. What we can see of the convection cells at the
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surface (basically, the top of them) is called a granule. With the huge surface area of the
Sun, there are millions of granules present. Their lifetime is, on average, fairly short, only
a few minutes, before that cell’s contents are reorganized into a different cell.
Why does the Sun even have a surface at all? The Sun is gaseous and therefore would
have no solid surface like a planet or a moon. What defines the surface? Recall that the
Sun’s energy transfer involved the absorption of a gamma ray and its subsequent reradiation. Heated materials can also excite electrons in an atom and when these excited
electrons get rid of their energy and drop back to its unexcited (ground) state, that “lost”
energy must be given off as light.
From the spectroscopy chapter, you know that there are many different forms of light, all
contained on the electromagnetic spectrum. So we can see the Sun because the Sun emits
visible-wavelength light, which means the Sun is hot enough at its surface to emit, as
visible light, the energy the atoms’ electrons are losing.
Does this mean that just above the surface, the temperature is not hot enough to excite the
electrons to a high enough level that they will emit visible light when they lose energy?
Actually, no. In fact, this lower part of the transparent Sun’s atmosphere, called the
chromosphere, is even hotter than the photosphere – roughly 20,000 K. So there should
be plenty of thermal energy to excite electrons.
The solution to this problem lies in the density of the Sun’s material. Like all bodies, the
Sun is densest at its core, and less dense towards the surface. This density gradient
continues beyond the photosphere; the chromosphere is less dense than the photosphere.
The crucial point about this density difference is that the Sun’s material is dense enough
in the photosphere (and in the convective zone in general) to be opaque to photons – in
other words, the photons are continually absorbed and re-emitted like they were in the
radiative zone and thus a single photon does not travel very far. In the less dense
chromosphere, there are not enough atoms packed in this layer to absorb the photons – so
they continue unimpeded into the upper layers of the solar atmosphere and interplanetary
space, and eventually to us.
This doesn’t mean that the photons don’t ever get absorbed and re-emitted in the
chromosphere; we can tell they do because there is some light coming from the
chromosphere, and from that, we can tell the composition of the chromosphere. Also,
much of the re-emitted photons from the chromosphere atoms corresponds to light we
cannot see – typically infrared light. We can tell this because some solar satellites are
equipped with infrared cameras that show the light coming from the chromosphere quite
well. The chromosphere is also beautifully imaged using the “hydrogen alpha” emission,
which turns out to be a deep red color.
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The chromosphere extends up to
about 10,000 km above the
photosphere. In comparison, the
Earth is about 12,700 km. The
chromosphere contains, at any
given instant, about 100,000
spicules, which are visible rising
columns of gas, lasting about five
minutes. There are other
structures as well, related to the
Sun’s magnetic field interacting
with these rising columns of
ionized gas and plasma:
filaments, prominences and
sunspots.
Above the chromosphere is an
even hotter layer of the solar
atmosphere called the corona. The corona has a temperature of about a million K and
extends a significant portion of the distance to Mercury’s orbit. The density of particles
here is so low (about a trillionth that of the photosphere which differs little from the
density of particles in interplanetary space. With a temperature of a million degrees, what
particles exist in the corona are travelling quite fast, so photons are not readily absorbed
and re-emitted by any of the atoms or ions there. So even though the corona is the largest
layer in volume of in all the Sun, it produces only one-millionth of the visible light the
Sun produces. In fact, to image the corona, you either need to view the Sun during a total
eclipse or using a coronagraph (which blocks out the photosphere artificially).
Why the corona is so hot is still debated, but the two principal mechanisms are induction
heating by the magnetic field of the Sun, and magnetohydrodynamic (MHD) waves
arising from the photosphere and chromosphere.
The corona, like the chromosphere, contains various structures: perhaps the most
dramatic of these are coronal mass ejections (CME), where a chunk of the plasma in the
corona detaches from the Sun and is directed away from the Sun. Therefore, the ejected
material is mostly electrons and protons. Because of the sheer amount of material, a CME
that is aimed at the Earth can have some far-reaching consequences, such as power
outages (as electrons are injected directly into power grids as extra current) and satellite
failures (as communications with satellites is completely disrupted, not allowing them to
fire their rockets to stay in orbit).
The Sun’s magnetosphere powers and directs the outbursts of solar activity
Before we get too far away from the Sun, consider where all this energy to drive CMEs
and similar phenomena comes from. The Sun, of course, radiates an immense amount of
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energy, but if it were just a matter of the heat of the core going through the various layers
of the Sun, there is nothing to suggest that there must be CMEs.
As implied in the previous section, much of the power and direction of solar outburst
phenomena arises from the Sun’s magnetic field, which itself is generated by the motion
of the electrically conductive fluid plasma in its convective zone. The Sun’s magnetic
field, like the Earth’s, varies over time due to its origin as random motions of its
convective cells and the Sun’s own rotation. Unlike any planetary magnetic field, though,
the Sun’s magnetic field varies in a completely predictable way over time. Every eleven
years, the Sun switches from having a standard north pole to south pole oriented
magnetic field (following the lines of longitude) to an east-west oriented magnetic field
(as if the lines are following lines of latitude) in the middle of the cycle, then back to a
north-south oriented magnetic field.
When there is a greater east-west
character to the Sun’s magnetic field
than a north-south character, the
magnetic force lines break through the
photosphere, where they create zones of
downwellling – that is, areas where the
cooler solar material sinks back into the
convective zone. These areas often
contain sunspots. These areas are also
places where the Sun contains active
features, such as flares, prominences
(projections of plasma in loop shapes
into the solar atmosphere, especially
near the edge of the Sun’s disk) and
filaments (dark lines across the
photosphere, which are prominences
seen from above).
The eleven-year cycle of the magnetic field of the Sun (and sunspot number) seems to be
related to the resilience of the solar magnetic field to distortion. The magnetic field lines
can be distorted by rotation until they are nearly east-west oriented instead of north-south,
but this distortion cannot carry on forever; the magnetic field seems to “snap” back into
the original north-south orientation.
Every eleven years, there is a maximum in solar activity – sunspot number, coronal mass
ejections and other photosphere disturbances all reach a maximum value. By contrast, at
solar minima, even the number of sunspots may drop to zero. These solar minima
coincide with years of cooler temperatures on Earth. For instance, during the period 1645
to 1715, which is several solar cycles in length, very few sunspots were recorded. This
period, known as the Maunder minimum for a couple who studied the behavior of
sunspots, also coincided with a period of significantly cooler temperatures worldwide
known as the Little Ice Age, though there is some controversy about its actual timing.
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Thus, the connection between solar activity minima and cooler temperatures on Earth is
not well understood.
Sunspots are areas of intense magnetic field activity
Sunspots form in those areas of the photosphere that are pierced by the Sun’s magnetic
field lines. Because they are slightly cooler (4200 K) than the surrounding photosphere
(6000 K), they appear darker than the
surrounding solar surface, but they are
hot enough to glow on their own. The
appearance of the sunspot is similar to
the top of a rising column of heated
water in a pan; the darkest most active
area is in the center of the sunspot (the
umbra), and it is surrounded by a
slightly less dark penumbra.
Sunspots form in pairs because any
magnetic field line that pierces the
photosphere from the inside out, must return to the inside and pierce the photosphere
again. All magnetic field lines emanate from the convective zone, and they all must
return to the convective zone.
The lifetime of an individual sunspot is a few days to a few months. After a solar
minimum, sunspots start appearing a relatively high (north and south) latitudes on the
Sun. As old sunspot pairs dissipate, new pairs form closer to the equator as the cycle
nears its maximum.
The graph above, sometimes called a “butterfly” graph due to its resemblance to the
insect, shows the date (year) along the horizontal axis and the position of sunspot groups
along the vertical axis. The solar minima are noticeable as there are almost no sunspots
anywhere on the Sun during those years (for example, 1890). As the cycle continues,
sunspots appear near 30°N and 30°S, then at lower latitudes (nearer the equator) as the
solar maximum is approached, before resetting at the next solar minimum.
The solar wind is an extension of the Sun into interplanetary space
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Less dramatic than sunspots and coronal mass ejections, but basically the same
phenomenon, is the solar wind, a steady stream of plasma directed away from the Sun in
all directions. There are two types of solar wind: the slow solar wind, which is made of
coronal material and is ejected as streamers away from the Sun’s equator, and the fast
solar wind, which is made of photospheric material and is ejected through coronal holes,
which occur near the Sun’s poles most often. The fast solar wind, true to its name, is
about twice as fast as the slow solar wind, but they both travel at hundreds of kilometers
per second.
According to a set of principles called Maxwell’s equations (see “Maxwell’s equations”
box), all moving charged particles, be they electrons, protons or ions, generate a magnetic
field. The energy to generate the magnetic field lies in the kinetic energy of the moving
particles themselves, meaning that the magnetic field is generated by slowing the charged
particles. Since the speed of the particles originated from the Sun’s processes, it is not a
stretch to say that the charged particles’ magnetic fields are a form of solar energy.
Either type of solar wind can be visualized in a couple of ways. The most dramatic is
viewing a comet’s tail. Comets are icy solar system bodies that orbit the Sun. As they
approach their perihelion, some of the ice begins to sublimate and form a gas, forming a
gaseous coma around the comet’s nucleus. Under the influence of relatively strong solar
winds near the Sun, some of the gas particles ionize. Directed by the magnetic fields
created by the solar wind, the cometary ions are directed away from the Sun, forming an
ion tail. The tail is visible because the ions contain excited electrons that emit light when
they lose their excitation. Because the ion tail is always directed away from the Sun, after
the comet’s perihelion and the comet moves away from the Sun, its ion tail will precede
the comet’s coma and nucleus as it moves.
The magnetosphere of a planet deflects the solar wind
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The other way to view the solar wind is when it encounters a planetary magnetic field,
which is generated due to the motion of an electrically conductive fluid. In the case of
Jupiter, it has a liquid metallic hydrogen layer whose motion generates the extremely
powerful magnetic field of Jupiter.
The Earth generates a magnetic field from motions of molten iron in its outer core. This
is the same magnetic field that allow compasses to point north. The Earth’s magnetic
field extends many Earth radii distance from the surface. Since Maxwell’s equations also
state that a charged particle will be influenced by a magnetic field, charged particles will
interact with the Earth’s magnetic field long before they can reach the Earth’s surface.
The region of the outer atmosphere where the Earth’s own magnetic field diverts the flow
of charged particles in the solar wind is called the magnetosphere.
An illustration of how good the magnetosphere is at diverting the solar wind particles is
the magnetopause, which is located a few hundred kilometers away from the surface (in
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fact, the International Space Station flies under this zone). At the magnetopause, the solar
wind speed has dropped from a few hundred kilometers per second to zero.
The Earth’s magnetosphere is not symmetrical. It is shaped like a teardrop, with the long
end pointing away from the Sun. There are a lot of closed magnetic field lines in the
magnetosphere, so many of the charged particles in the solar wind are trapped, moving
back and forth along these magnetic field lines. One consequence of this behavior is that
those high energy electrons and protons do not reach the surface of the Earth, where life
would be affected.
The aurora is a phenomenon which is related to the trapped electrons and protons in the
magnetosphere. Interestingly, a particularly energetic phase of a solar wind will give
charged particles enough energy that they
will be diverted to the tail of the
magnetosphere (the pointy end of the
teardrop) where they will still have enough
energy to follow magnetic field lines back
towards the Earth. As they encounter
nitrogen and oxygen molecules in the
upper atmosphere, the particles will excite
electrons in those molecules. As the
electrons relax, they emit light, as they
must. Oxygen molecules emit red and
green; nitrogen molecules emit pink and
blue/violet. And because the charged
particles are coming from the “back” side of the Earth, auroras are particularly active on
the night side of the Earth.
The Earth is by no means the only planet with a magnetosphere; in fact, auroras have
been observed on all planets except Mercury. In fact, some moons have auroras: Io,
Europa and Ganymede (three of Jupiter’s Galilean satellites) display them.
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The Van Allen radiation belts contain significant ionizing radiation
Another expression of
Earth’s magnetosphere are
the Van Allen radiation
belts, named after James
Van Allen, the University
of Iowa space scientist
who discovered them in
1958, by placing a Geiger
counter on one of the first
orbiting satellites,
Explorer I. There are two
belts, the inner radiation
belt which is 650 to 6500
km above the surface, and
the outer radiation belt
which is 13000 to 50000
km. Both are toroid
(doughnut) in shape, with the “hole” centered over the poles.
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The main difference between the two belts is the source of the radiation. The inner belt is
populated by protons generated by the collision of cosmic ray particles and upper
atmosphere atoms. These protons are stabilized by a process called magnetic mirroring
and can last for years, moving back and forth between the magnetic poles, following
Earth’s magnetic field lines.
The outer belt is populated by solar wind electrons funneled in from the tail of the
magnetosphere. These electrons have much less energy than the protons of the inner belt.
Both belts are significant hazards to spacecraft electronics, causing power surges and
other disruptions. Though both belts have measurable amounts of ionizing radiation
(which can affect living tissue), astronauts could live for several months within each belt
without receiving a lethal dose of radiation. To avoid any such problems, the Apollo
missions, the only mission with humans to reach the height of the Van Allen belts, skirted
the edges of the belts on their way to the Moon.
Geomagnetic storms can have worldwide effects
A geomagnetic storm occurs when there is a measureable increase in the amount of
charged particles, typically originating in the Sun. In the same way that precipitation on
Earth’s surface can measure the severity of an atmospheric storm, the amount of charged
particles is used to measure the severity of a geomagnetic storm.
For instance, during the solar minima, storms may be caused by the fast solar wind (from
the coronal holes) overtaking the slow solar wind and creating a shock wave of particles
that overloads the magnetosphere.
During the solar maxima, storms may be caused by the direct injection of a huge number
of charged particles from a coronal mass ejection.
There are a couple of reasons why we worry about geomagnetic storms: the effect on
humans and the effect on our technological infrastructure.
Clearly, if you are on the Earth’s surface, the Earth’s magnetosphere does a near-perfect
job of shielding you from any amount of solar or cosmic charged particles, mostly be
being dense enough to interact with, and thus lessen the energy of, these charged particles
so that they cannot penetrate further into the atmosphere. However, if you are in space, or
even on a plane, you lose this shielding to varying degrees. So the exposure of astronauts,
airplane crew and passengers to charged particles must be measured. On a cross-country
flight from Los Angeles to New York, for instance, you are exposed to as much ionizing
radiation as you would if you had gotten two chest X-rays or eaten 400 bananas.
However, the risks to our civilization’s technological infrastructure, specifically the
electrical distribution grid, is much more significant. For instance, a solar flare in 1989
caused an overload in the electrical grid of Quebec, causing a blackout that affected nine
million people.
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The modern convention used to describe solar flare strength (and therefore possible
effects on Earth) is the ABCMX scale, which rates flares based on the intensity of X-rays
in the 0.1 nm to 0.8 nm range. A is the least strong classification and X is the strongest,
as shown in the table below.
Solar
flare
class
A
B
C
M
X
Intensity of maximum
flare activity (W/m2)
10–8
10–7
10–6
10–5
10–4
The letter designations are often followed by a number, which gives a more precise flare
strength. For instance, a M2.2 flare, such as the one that occurred March 12, 2015,
generates X-rays with a maximum intensity of 2.2 × 10–5 W/m2 – strong enough to cause
an hour-long blackout of some radio frequencies. Thus, 9.9 is the largest number
associated with the various letter classes. Since there isn’t a classification higher than X,
X-class flares can have numbers much greater than 9.9. Just before noon on September 1,
1859, solar astronomer Richard Carrington was observing the sun in his private
observatory in England, when he noticed two blindingly white “beads” appear over the
dark sunspots he was studying. By the time he found a witness to corroborate his
observation, the “beads” were gone. By dawn the next morning, auroras were sighted all
over the world, even as close to the equator as Hawaii and El Salvador. Modern estimates
of the so-called Carrington flare place it at X42 – translated to an intensity of 42 × 10–4
W/m2, or nearly 200 times the intensity of the 2015 flare.
How about some space travel?
Because the slow solar wind is more or less constant over time, various researchers have
suggested intercepting the stream of flowing particles with a solar sail, and using that as
a propulsion system,
similar to the wind
propelling sailing ships
on oceans.
The solar wind could
be used to steer, and
even to move towards
the Sun.
The most recent
successful version of
this type of propulsion
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technology to be launched was NASA’s NanoSail-D2, a ten square meter, four kilogram
sail that was deployed in 2010, and stayed in orbit for 240 days. The sail material itself
was a thin piece of reflective ultra-thin polymer plastic, and it accomplished the goal of
deploying without tangling itself up, a problem that earlier solar sails had encountered.
Of course, this sail was deployed in low Earth orbit, so it didn’t really use the solar wind
for propulsion; there has yet to be a true deep-space test of solar sails.
In the next chapter, you’ll see a couple reasons why we can’t simply test propulsion
technologies like solar sails in interplanetary space.
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