Transcript Distances in Space

Distances in Space
Distance Units - the AU
An Astronomical Unit (AU) is defined as the average
distance between the Earth and Sun.
1 AU = 92 955 807 miles or 149 597 871 km
AU’s are used to measure distances within the solar system.
There are approximately 63,240 AU’s in one light-year (LY).
Distance Units - the Light Year
One light year (LY) is the distance light travels in one year.
Light years are not used to measure time.
A light year is a convenient unit to use for measuring
distances within the galaxy, especially the distances to
nearby stars.
One LY = 5.88 trillion miles or 9.46 trillion km.
Distance Units - Parsecs
A parsec is a shortened term for “parallax second of
arc”. A star at a distance of one parsec shows a parallax
angle of 1 second of 1 degree (1/3600 of a degree).
This translates to a distance of 3.26 LY.
Distances in parsecs are not commonly quoted, but
diatances in kiloparsecs (kpc), megaparsecs (Mpc) or
gigaparsecs (Gpc) are often used, especially when
referring to distances to faraway galaxies.
1 kpc = 1000 pc = 3260 LY
1Mpc = 1 million pc = 3.26 million LY
1Gpc = 1 billion pc = 3.26 billion LY
The Distance Ladder
Calculating distances in
space is much like
climbing a ladder:
Each step higher relies on
the previous step below,
and the higher you go,
the more uncertain you
are about your stability.
The Distance Ladder
Calculating distances in space is much like climbing a
ladder:
Each step higher relies on the previous step below,
and the higher you go, the more uncertain you are
about your stability .
Knowing the distance to one object often serves as a
stepping stone to determining the distance to another.
The Distance Ladder
Calculating distances in space is much like climbing a
ladder:
Each step higher relies on the previous step below,
and the higher you go, the more uncertain you are
about your stability .
Knowing the distance to one object often serves as a
stepping stone to determining the distance to another.
The most distant measurements are those with the
highest degrees of uncertainties.
The Distance Ladder
Preliminary Measurements
The early Greeks had to first determine the diameters
of the Earth and Moon before they could calculate a
distance to the Moon.
Eratosthenes performed the now famous experiment
between Syene and Alexandria to determine the Earth’s
circumference, and thus its diameter. For a review of
the details of the experiment, check out this link:
http://dev.physicslab.org/Document.aspx?doctype=2&f
ilename=IntroductoryMathematics_EarthCircumferenc
e.xml
Earth Circumference
A short excerpt of the Eratosthenes experiment from
the Cosmos with Carl Sagan explains this idea:
https://www.youtube.com/watch?v=G8cbIWMv0rI
A teachers guide to replicate the activity is here:
http://www.physics2005.org/projects/eratosthenes/Te
achersGuide.pdf
In order to really do this, you must find another school
or observer that is along your line of longitude, and at a
some distance away.
Size of the Moon
Knowing the size of the Earth, Aristotle used observations
from a lunar eclipse to get an estimate of the Moon’s
angular size. An interactive of this activity is at this link,
but you can easily turn it into a pencil and paper activity:
http://dev.physicslab.org/Document.aspx?doctype=2&file
name=IntroductoryMathematics_SizeMoon.xml
Distance to the Moon
Once Aristotle knew the Moon’s diameter, he could calculate a
distance to the Moon by using the property of similar triangles.
(You may want to first ask students to experiment with different
sized coins to see how far they have to move the coin away from
their eye to get it to match the size of the Moon in the sky. (But
even with a dime, this can be challenging.)
The following activity suggests using a thumbnail, but using (the
width of) a pencil works better:
http://dev.physicslab.org/Document.aspx?doctype=2&filename=
IntroductoryMathematics_EarthMoonDistance.xml
Distance to the Moon
Once Aristotle knew the Moon’s diameter, he could
calculate a distance to the Moon by using the property of
similar triangles
Distance to the Moon
The Apollo XI. astronauts left a mirror on the lunar
surface so that laser beams from Earth could be
reflected back. The measurement of the incoming light
gives a highly accurate measurement of the Moon’s
distance, which varies from 363,104 km – 406,696 km,
depending on where it is in its orbit.
The link to the full story is here:
http://science.nasa.gov/science-news/science-atnasa/2004/21jul_llr/
Distance to the Moon
The lunar laser ranging experiment is still operational
and can determine the distance between Earth and
Moon to a few centimeters.
The McDonald Observatory in Texas
operates the laser that bounces off
the Moon.
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Distance to the Sun – Aristarchus
http://galileoandeinstein.physics.virginia.edu/lectures/gkastr1.html
When the Moon appears to be exactly half lit, the line from the Moon
to the Sun must be exactly perpendicular to the line from the Moon to
an observer on Earth. (See the figure below). So, if the quarter moon is
observed during the day, one could measure the angle between the
Moon and the Sun in the sky (the angle α in the figure). You can then
construct a long thin triangle, with its baseline the Earth-Moon line,
having an angle of 90 degrees at one end and α at the other, and so
find the ratio of the Sun’s
distance to the Moon’s
distance.
This may be solved using either
trigonometry or a scale drawing.
Distance to the Sun – scale drawing
The Earth-Moon distance is 384 400 km (adjacent side).
Use a scale of 1 cm = 100 000 km for your drawing.
Use a protractor to create angles of 90° and 89.8° on the drawing.
Draw out the opposite and hypotenuse sides until they intersect. From
that point, draw a straight line back to the adjacent side to represent
the Earth-Sun distance.
Measure this in cm and convert it
to km to get the distance from
Earth to the Sun.
Distance to the Sun – more to come
There wasn’t any good way to determine a more accurate
distance to the Sun until after Kepler had derived his three laws
of planetary motion, so we will come back to this idea once
Kepler’s Laws have been examined.
Sneak preview:
In 1677, English astronomer Edmond Halley proposed a method
for calculating the Earth’s distance from the Sun by using the
transit of Venus.
Halley died 19 years before his method could be attempted and
proven successful (during the 1761 transit of Venus).
Kepler’s Third Law
Johannes Kepler derived three laws of planetary motion, which
revolutionized how astronomers understood the solar system.
His third law, published in 1618, gave a way to measure the
relative distance of a planet from the Sun:
P2 = A3
where P = planet's orbital period in years
and A = the planet’s average distance from the Sun, in AU
Example: Venus orbits the Sun in .615 Earth years.
What is its distance from the Sun?
.6152 = A3
A= .723 AU
A Transit of Venus
A Venus transit is when Venus
crosses exactly between Earth and
the Sun. The silhouette of the planet
may be viewed moving across the
solar disc, like a miniature eclipse.
This alignment is rare, coming in
pairs that are eight years apart but
separated by over a century. The
most recent transit of Venus was in
June of 2012. The next transit of
Venus pair occurs in December 2117
and 2125.
Distance to the Sun - transit of Venus - 1
http://www.exploratorium.edu/venus/question4.html
Imagine two different people, at two different locations on Earth, viewing
the transit of Venus.
The angle between the two different paths measured from Earth is angle E:
Distance to the Sun - transit of Venus - 2
From Kepler's third law, we know the relative distances of all the planets
from the Sun. Case in point: We know that Venus's distance from the Sun is
0.72 times the Earth's distance from the Sun.
This distance relationship also tells us angle V, the angle between the two
paths as seen from Venus: angle V is angle E divided by 0.72. (This is true
only for small angles, which these are.)
Distance to the Sun - transit of Venus - 3
In addition to angle V, we also need to know the distance
between the two observers on Earth, at points A and B. Call
this distance dA-B.
For small angles, tan (1/2 A) = 1/2 tan A, so the distance
between Earth and Venus is equal to dA-B / tan V. Note that
this gives only the distance to Venus, not to the Sun.
Distance to the Sun - transit of Venus - 4
If we have the distance between Earth and Venus, then the
distance from Earth to the Sun may be calculated by a simple
ratio:
Venus
distance
.72 AU
107,000,000 km
= 1.0 AU
x km
Earth
distance
*X= 148,600,000 km
*This was the distance between Earth and Sun at the time of the
2012 transit.
The average distance between Earth and Sun is 149,597,871 km.
Distance from Radar
In modern times, radar has
been used to measure
distance to the inner planets
and asteroids, as well as
comets.
The Arecibo radio telescope
in Puerto Rico is one of the
instruments that is used for
these studies.
The Arecibo telescope has
the capability to transmit a
high powered-beam of radio
waves, as well as to detect
their returning echoes.
Distance from Radar
Although radar may be
used to determine
distances, its primary
function is to create maps
of planetary surfaces and
to determine the shapes
of asteroids and comet
nuclei.
Radar image of the Moon’s south
polar region.
Parallax
The parallax technique is used to measure the distances to
nearby stars, out to about 500 LY. (The more distant stars have
such small motions that an accurate parallax angle cannot be
determined for them.)
The baseline of Earth’s orbit is used to measure the parallax
angle. This is the angular amount of perceived shift of the
nearby star against the backdrop of the distant stars.
Parallax and Parsces
There is a simple relationship between a star's distance and its
parallax angle:
d = 1/p
The distance d is measured in parsecs and the parallax angle p is
measured in arcseconds.
This relationship is why many astronomers prefer to measure
distances in parsecs.
Star Magnitudes
Many of the subsequent distance techniques rely on knowing star magnitudes.
There are two types: Apparent (m) and Absolute (M)
Apparent magnitude is how bright the star appears from Earth.
Absolute magnitude is how bright the star would appear if placed at a distance of
10 parsecs.
So if m < M, the star would be closer than 10 parsecs.
Both magnitude scales work the same way: the more negative the number, the
brighter the star.
The Sun has an absolute magnitude of 4.74 and an apparent magnitude of -26.7.
Spectroscopic Parallax - definition
Spectroscopic parallax is a sort of misnomer, because
there is no measurement of a shift in position, as there
is in (trigonometric) parallax.
The use of the word ‘parallax’ in this instance refers to
using a star’s spectrum to find its distance.
This technique can be applied to any star, assuming
that its apparent magnitude has been measured and a
spectrum has been recorded.
Steps for Spectroscopic Parallax
1. The pattern of spectral lines assigns a spectral class to the
star, giving it a horizontal position on an H-R diagram.
Steps for Spectroscopic Parallax
2. The thickness of spectral lines is used to assign a luminosity
class to the star, giving it a vertical position on an H-R diagram.
Steps for Spectroscopic Parallax
3. From its location on H-R
Diagram, an absolute
magnitude can be
estimated.
This links to an applet that may
be used to demonstrate these
steps:
http://astro.unl.edu/naap/distance/anim
ations/spectroParallax.html
Steps for Spectroscopic Parallax
4. The distance modulus equation can then be used to calculate
the distance to a star:
d = 10(m - M + 5)/5
where d is the distance in parsecs,
m is the apparent magnitude of the star and
M is the absolute magnitude of the star
Calculating spectroscopic parallax is one of the most important
applications of the H-R Diagram.
The Distance Modulus
In addition to the
equation, the distance
modulus can be easily
applied through the use
of a nomogram.
Draw a straight line
connecting the two
magnitudes of a star (M
and m).
The distance may be read
from where the line
crosses the middle axis.
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Standard Candles
Some objects in space have fixed and known absolute
magnitudes, and can be used to determine distances by using
the distance modulus. Since their intrinsic brightness is known,
their apparent magnitude is directly an effect of distance.
Think of it like a 100W light bulb: As the distance between you
and the light bulb doubles, the brightness drops off according to
the inverse-square law:
Brightness = 1/d2
Doubling the distance makes the light
bulb appear ¼ of its original brightness.
Examples of Standard Candles
Here are some celestial objects that are used as standard candles:
 Cepheid and RR Lyrae variable stars
 Type Ia supernovae
 Spiral galaxies with measured rotational velocities (Tully-Fisher
relationship)
 Gravitationally-lensed quasars
There are other standard candles that are measured using radio or
x-rays wavelengths.
Cepheid variable stars
Cepheid variable stars may be used to measure the distance
to a star cluster or nearby galaxy. This technique is good
out to distances of 60 million LY.
They obey a well-defined period-luminosity
relationship, discovered in the early 1900’s
by the Harvard astronomer Henrietta Leavitt.
This is now known as the Leavitt Law.
Other types of variable stars were subsequently discovered
with similar patterns, such as RR Lyrae variables.
Cepheid variable stars
To find the distance to a Cepheid variable star, one
must first observe its brightness fluctuations over time
and plot a light curve. Then its period and average
magnitude may be calculated.
This plot shows a
star with a period of
5 days and an average
apparent magnitude
of 15.6:
Cepheid variable stars
Cepheid variable stars obey a well-defined periodluminosity relationship:
The luminosity can be converted to an absolute
magnitude, and the distance
modulus equation is then
used to calculate distance.
Star Cluster Fitting
Another way to use the H-R Diagram to calculate
distance is to “fit” the main sequence of a star cluster
to a model or standard star cluster. In this method, it is
assumed that all the stars of a cluster are at the same
distance away, so when it is fitted, the distance to the
cluster is determined.
The model cluster contains stars whose distances are
known from parallax and therefore have defined
absolute magnitudes.
Star Cluster Fitting
In this example, a star cluster is plotted (grey dots)
against the model track of a known star cluster (red).
The vertical distance is adjusted between the two main
sequences so that
they overlap.
The stars in this
cluster show dimmer
apparent magnitudes,
so this cluster must be
farther away than the
model cluster.
Star Cluster Fitting
The vertical distance is adjusted between the two main
sequences so that they overlap. The difference (Δm) may be
used in the distance modulus equation as a substitute for (M-m):
d = 10((Δ m) + 5)/5
where d is the distance in parsecs between the two clusters
In the previous example, the Δm was 5.4 magnitudes, so
d = 10((5.4) + 5)/5 = 102.08 = 120 parsecs
If the model cluster had a known distance of 80 pc, then the
fitted cluster distance is 80 + 120 = 200 pc.
Star Cluster Fitting
Here’s an interactive activity:
http://astro.unl.edu/naap/distance/animations/clusterFittingExplorer
.html
Supernova Type Ia light curves
The light curves from Type Ia supernovae may be used as distance indicators to
distant galaxies by using a variation of the distance modulus equation:
d = 10 (m-M+5)/5
Where distance (d) is the distance in parsecs,
m is the peak apparent magnitude of the
supernova and M is the average absolute
magnitude of all “standard candle” Type Ia
supernovae, equal to -19.3.
Supernova Type Ia light curves
As more data from SN light curves accumulated, it was realized that not all Type
Ia SN behaved exactly alike. Below is a refinement technique introduced by
Mark Phillips in 1993.
There’s a correlation between the peak brightness of the supernova and the
rate at which its brightness declines over a 15 day period after maximum light
in the B (blue filter) band. It is known as the luminosity-decline relation:
Brighter SN fade more slowly, and fainter ones more rapidly.
The change in blue magnitude in the 15 days following maximum light is Δ m15.
Galaxy redshifts
Most galaxies show a shift in
their spectral lines towards
longer (redder) wavelengths.
This is primarily due to the
expansion of the universe.
By measuring the amount of
the shift, the line-of-sight
velocity and distance to the
galaxy may be calculated
using Hubble’s law.
Measuring galaxy redshifts
The first step in the procedure is to measure the
displacement of certain spectral lines from their original
(rest) position.
The rest wavelength is the laboratory position of the line. The
displaced position will be called the observed wavelength.
The redshift (Z) is calculated by this equation:
Z = (observed wavelength/ rest wavelength) - 1
Note: Z should be a positive number
Galaxy velocity from redshift
Once Z is known, the velocity (v) in km/sec may be calculated by
using this equation below:
c (the speed of light) = 3.0 x 105 km/sec
v = (Z + 1)2 -1
c = (Z + 1)2 +1
Galaxy distance from velocity
Hubble’s law states that a galaxy’s distance is
proportional to its velocity, by this relationship:
Distance (in Mpc) = (cz/H) * ((1 + 0.5 z)/(1 + z))
Where H is the Hubble constant. Present values
place this number between 67.3 - 73.8
km/sec/Mpc
Distances to Quasars
The distance to a quasar may be calculated in the same
manner as that distance to a galaxy. There is one important
difference: Quasar spectra (left) are marked by numerous
emission lines, while galaxy spectra (right) contain primarily
absorption lines.
Seeing far away Galaxies: Gravitational Lensing