The KStars Handbook

Now that we have the time and location set, let us have a look around. You can pan the display using the arrow keys. If you hold down the Shift key before panning, the scrolling speed is doubled. The display can also be panned by clicking and dragging with the mouse. Note that while the display is scrolling, not all objects are displayed. This is done to cut down on the CPU load of recomputing object positions, which makes the scrolling smoother (you can configure what gets hidden while scrolling in the Configure KStars window). There are seven ways to change the magnification (or Zoom level) of the display:

Notice that as you zoom in, you can see fainter stars than at lower zoom settings.

Zoom out until you can see a green curve; this represents your local horizon. If you have not adjusted the default KStars configuration, the display will be solid green below the horizon, representing the solid ground of the Earth. There is also a white curve, which represents the celestial equator, and a tan curve, which represents the Ecliptic, the path that the Sun appears to follow across the sky over the course of a year. The Sun is always found somewhere along the Ecliptic, and the planets are never far from it.

KStars displays thousands of celestial objects: stars, planets, comets, asteroids, clusters, nebulae and galaxies. You can interact with displayed objects to perform actions on them or obtain more information about them. Clicking on an object will identify it in the status bar, and simply hovering the mouse cursor on an object will label it temporarily in the map. Double-clicking will re-centre the display on the object and begin tracking it (so that it will remain centred as time passes). Right clicking an object opens the object's popup menu, which provides more options.

The Popup Menu

Here is an example of the right click popup menu, for the Orion Nebula:

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The appearance of the popup menu depends somewhat on the kind of object you right-click on, but the basic structure is listed below. You can get more detailed information about the popup menu.

The top section contains information labels (which are not selectable). The top one to three labels display the object's name(s) and object type. The next three labels show the object's rise, transit and set times. If the rise and set times say "circumpolar", it means that the object is always above the horizon for the present location.

The middle section contains items for performing actions on the object, such as Center and Track, Angular Distance To... (which enters the Angular Distance Mode, allowing you to measure the angle between objects in the sky), Details... (which opens the object's Object Details window), Attach Label and Add/Remove Trail (only available for Solar System bodies).

The bottom section contains links to images and/or informative webpages about the selected object. If you know of an additional URL with information or an image of the object, you can add a custom link to the object's popup menu using the Add Link... item.

Finding Objects

You can search for named objects using the Find Object tool, which can be opened by clicking on the search icon in the toolbar, by selecting Find Object... from the Pointing menu, or by pressing Ctrl+F. The Find Object window is shown below:

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The window contains a list of all the named objects that KStars is aware of. Many of the objects only have a numeric catalogue name (for example, NGC 3077), but some objects have a common name as well (for example, Whirlpool Galaxy). You can filter the list by name and by object type. To filter by name, enter a string in the edit box at the top of the window; the list will then only contain names which start with that string. To filter by type, select a type from the combo box at the bottom of the window.

To centre the display on an object, highlight the desired object in the list, and press Ok. Note that if the object is below the horizon, the program will warn you that you may not see anything except the ground (you can make the ground invisible in the Display Options window, or by pressing the Ground button in the View toolbar).

KStars will automatically begin tracking on an object whenever one is centred in the display, either by using the Find Object window, by double-clicking on it, or by selecting Center and Track from its right-click popup menu. You can disengage tracking by panning the display, pressing the Lock icon in the Main toolbar, or selecting Track Object from the Pointing menu.

Note

When tracking on a Solar System body, KStars will automatically attach an “orbit trail”, showing the path of the body across the sky. You will likely need to change the clock's timestep to a large value (such as “1 day”) to see the trail.

This concludes the tour of KStars, although we have only scratched the surface of the available features. KStars includes many useful astronomy tools, it can directly control your telescope, and it offers a wide variety of configuration and customization options. In addition, this Handbook includes the AstroInfo Project, a series of short, interlinked articles explaining some of the celestial and astrophysical concepts behind KStars.

When KStars starts up, the time is set to your computer's system clock, and the KStars clock is running to keep up with the real time. If you want to stop the clock, select Stop Clock from the Time menu, or simply click on the Pause icon in the toolbar. You can make the clock run slower or faster than normal, or even make it run backward, using the time-step spinbox in the toolbar. This spinbox has two sets of up/down buttons. The first one will step through all 83 available time steps, one by one. The second one will skip to the next higher (or lower) unit of time, which allows you to make large timestep changes more quickly.

You can set the time and date by selecting Set Time... from the Time menu, or by pressing the time icon in the toolbar. The Set Time window uses a standard KDE Date Picker widget, coupled with three spinboxes for setting the hours, minutes and seconds. If you want to re-synchronise the simulation clock back to the current CPU time, just select Set Time to Now from the Time menu.

The Configure KStars window allows you to modify a wide range of display options. You can access the window with the configure toolbar icon, or by selecting Configure KStars... from the Settings menu. The window is depicted below:

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The Configure KStars window is divided into five tabs: Catalogues, Guides, Solar System, Colours and Advanced.

In the Catalogues tab, you determine which object catalogues are displayed in the map. The Stars section also allows you to set the “faint magnitude limit” for stars, and the magnitude limit for displaying the names and/or magnitudes of stars. Below the stars section, the Deep-Sky Objects section controls the display of several non-stellar object catalogues. By default, the list includes the Messier, NGC and IC catalogues. You can add your own custom object catalogues by pressing the Add Custom Catalog button. For detailed instructions on preparing a catalogue data file, see the README.customize file that ships with KStars.

In the Solar System tab, you can specify whether the Sun, Moon, planets, comets and asteroids are displayed, and whether the major bodies are drawn as coloured circles or actual images. You can also toggle whether solar system bodies have name labels attached, and control how many of the comets and asteroids get name labels. There is an option to automatically attach a temporary “orbit trail” whenever a solar system body is tracked, and another to toggle whether the colour of the orbit trail fades into the background sky colour.

The Guides tab lets you toggle whether non-objects are displayed (i.e., constellation lines, constellation names, the Milky Way contour, the celestial equator, the ecliptic, the horizon line, and the opaque ground). You can also choose whether you would like to see Latin constellation names, IAU-standard three-letter abbreviations, or constellation names using your local language.

The Colours tab allows you to set the colour scheme, and to define custom colour schemes. The tab is split into two panels:

The left panel shows a list of all display items with adjustable colours. Click on any item to bring up a colour selection window to adjust its colour. Below the list is the Star Colour Mode selection box. By default, KStars draws stars with a realistic colour tint according to the spectral type of the star. However, you may also choose to draw the stars as solid white, black or red circles. If you are using the realistic star colours, you can set the saturation level of the star colours with the Star Colour Intensity spinbox.

The right panel lists the defined colour schemes. There are four predefined schemes: the Default scheme, Star Chart, which uses black stars on a white background, Night Vision, which uses only shades of red in order to protect dark-adapted vision, and Moonless Night, a more realistic, dark theme. Additionally, you can save the current colour settings as a custom scheme by clicking the Save Current Colors button. It will prompt you for a name for the new scheme, and then your scheme will appear in the list in all future KStars sessions. To remove a custom scheme, simply highlight it in the list, and press the Remove Colour Scheme button.

The Advanced Tab provides fine-grained control over the more subtle behaviours of KStars.

The Correct for atmospheric refraction checkbox controls whether the positions of objects are corrected for the effects of the atmosphere. Because the atmosphere is a spherical shell, light from outer space is “bent” as it passes through the atmosphere to our telescopes or eyes on the Earth's surface. The effect is largest for objects near the horizon, and actually changes the predicted rise or set times of objects by a few minutes. In fact, when you “see” a sunset, the Sun's actual position is already well below the horizon; atmospheric refraction makes it seem as if the Sun is still in the sky. Note that atmospheric refraction is never applied if you are using Equatorial coordinates.

The Use animating slewing checkbox controls how the display changes when a new focus position is selected in the map. By default, you will see the sky drift or “slew” to the new position; if you untick this option, then the display will instead “snap” immediately to the new focus position.

If the Attach label to centred object checkbox is selected, then a name label will automatically be attached to an object when it is being tracked by the program. The label will be removed when the object is no longer being tracked. Note that you can also manually attach a persistent name label to any object with its popup menu.

There are three situations when KStars must redraw the sky display very rapidly: when a new focus position is selected (and Use animated slewing is checked), when the sky is dragged with the mouse, and when the time step is large. In these situations, the positions of all objects must be recomputed as rapidly as possible, which can put a large load on the CPU. If the CPU cannot keep up with the demand, then the display will seem sluggish or jerky. To mitigate this, KStars will hide certain objects during these rapid-redraw situations, as long as the Hide objects while moving checkbox is selected. The timestep threshold above which objects will be hidden is determined by the Also hide if timescale greater than: timestep-spinbox. You can specify the objects that should be hidden in the Configure Hidden Objects group box.

There are several ways to modify the display to your liking.

A basic requirement for studying the heavens is determining where in the sky things are. To specify sky positions, astronomers have developed several coordinate systems. Each uses a coordinate grid projected on the Celestial Sphere, in analogy to the Geographic coordinate system used on the surface of the Earth. The coordinate systems differ only in their choice of the fundamental plane, which divides the sky into two equal hemispheres along a great circle. (the fundamental plane of the geographic system is the Earth's equator). Each coordinate system is named for its choice of fundamental plane.

The Celestial Equator is an imaginary great circle on the celestial sphere. The celestial equator is the fundamental plane of the Equatorial Coordinate System, so it is defined as the locus of points with Declination of zero degrees. It is also the projection of the Earth's equator onto the sky.

The Celestial Equator and the Ecliptic are set at an angle of 23.5 degrees in the sky. The points where they intersect are the Vernal and Autumnal Equinoxes.

The sky appears to drift overhead from east to west, completing a full circuit around the sky in 24 (Sidereal) hours. This phenomenon is due to the spinning of the Earth on its axis. The Earth's spin axis intersects the Celestial Sphere at two points. These points are the Celestial Poles. As the Earth spins; they remain fixed in the sky, and all other points seem to rotate around them. The celestial poles are also the poles of the Equatorial Coordinate System, meaning they have Declinations of +90 degrees and -90 degrees (for the North and South celestial poles, respectively).

The North Celestial Pole currently has nearly the same coordinates as the bright star Polaris (which is Latin for “Pole Star”). This makes Polaris useful for navigation: not only is it always above the North point of the horizon, but its Altitude angle is always (nearly) equal to the observer's Geographic Latitude (however, Polaris can only be seen from locations in the Northern hemisphere).

The fact that Polaris is near the pole is purely a coincidence. In fact, because of Precession, Polaris is only near the pole for a small fraction of the time.

The celestial sphere is an imaginary sphere of gigantic radius, centred on the Earth. All objects which can be seen in the sky can be thought of as lying on the surface of this sphere.

Of course, we know that the objects in the sky are not on the surface of a sphere centred on the Earth, so why bother with such a construct? Everything we see in the sky is so very far away, that their distances are impossible to gauge just by looking at them. Since their distances are indeterminate, you only need to know the direction toward the object to locate it in the sky. In this sense, the celestial sphere model is a very practical model for mapping the sky.

The directions toward various objects in the sky can be quantified by constructing a Celestial Coordinate System.

The ecliptic is an imaginary Great Circle on the Celestial Sphere along which the Sun appears to move over the course of a year. Of course, it is really the Earth's orbit around the Sun causing the change in the Sun's apparent direction. The ecliptic is inclined from the Celestial Equator by 23.5 degrees. The two points where the Ecliptic crosses the Celestial Equator are known as the Equinoxes.

Since our solar system is relatively flat, the orbits of the planets are also close to the plane of the ecliptic. In addition, the constellations of the zodiac are located along the ecliptic. This makes the ecliptic a very useful line of reference to anyone attempting to locate the planets or the constellations of the zodiac, since they all literally “follow the Sun”.

Because of the 23.5-degree tilt of the Ecliptic, the Altitude of the Sun at noon changes over the course of the year, as it follows the path of the Ecliptic across the sky. This causes the seasons. In the Summer, the Sun is high in the sky at noon, and it remains above the Horizon for more than twelve hours. Whereas, in the winter, the Sun is low in the sky at noon, and remains above the Horizon for less than twelve hours. In addition, sunlight is received at the Earth's surface at a more direct angle in the Summer, which means that a given area at the surface receives more energy per second in the Summer than in Winter. The differences in day duration and in energy received per unit area lead to the differences in temperature we experience in Summer and Winter.

Most people know the Vernal and Autumnal Equinoxes as calendar dates, signifying the beginning of the Northern hemisphere's Spring and Autumn, respectively. Did you know that the equinoxes are also positions in the sky?

The Celestial Equator and the Ecliptic are two Great Circles on the Celestial Sphere, set at an angle of 23.5 degrees. The two points where they intersect are called the Equinoxes. The Vernal Equinox has coordinates RA=0.0 hours, Dec=0.0 degrees. The Autumnal Equinox has coordinates RA=12.0 hours, Dec=0.0 degrees.

The Equinoxes are important for marking the seasons. Because they are on the Ecliptic, the Sun passes through each equinox every year. When the Sun passes through the Vernal Equinox (usually on March 21st), it crosses the Celestial Equator from South to North, signifying the end of Winter for the Northern hemisphere. Similarly, whenthe Sun passes through the Autumnal Equinox (usually on September 21st), it crosses the Celestial Equator from North to South, signifying the end of Winter for the Southern hemisphere.

Locations on Earth can be specified using a spherical coordinate system. The geographic (“earth-mapping”) coordinate system is aligned with the spin axis of the Earth. It defines two angles measured from the centre of the Earth. One angle, called the Latitude, measures the angle between any point and the Equator. The other angle, called the Longitude, measures the angle along the Equator from an arbitrary point on the Earth (Greenwich, England is the accepted zero-longitude point in most modern societies).

By combining these two angles, any location on Earth can be specified. For example, Baltimore, Maryland (USA) has a latitude of 39.3 degrees North, and a longitude of 76.6 degrees West. So, a vector drawn from the centre of the Earth to a point 39.3 degrees above the Equator and 76.6 degrees west of Greenwich, England will pass through Baltimore.

The Equator is obviously an important part of this coordinate system; it represents the zeropoint of the latitude angle, and the halfway point between the poles. The Equator is the Fundamental Plane of the geographic coordinate system. All Spherical Coordinate Systems define such a Fundamental Plane.

Lines of constant Latitude are called Parallels. They trace circles on the surface of the Earth, but the only parallel that is a Great Circle is the Equator (Latitude=0 degrees). Lines of constant Longitude are called Meridians. The Meridian passing through Greenwich is the Prime Meridian (longitude=0 degrees). Unlike Parallels, all Meridians are great circles, and Meridians are not parallel: they intersect at the north and south poles.

Consider a sphere, such as the Earth, or the Celestial Sphere. The intersection of any plane with the sphere will result in a circle on the surface of the sphere. If the plane happens to contain the centre of the sphere, the intersection circle is a Great Circle. Great circles are the largest circles that can be drawn on a sphere. Also, the shortest path between any two points on a sphere is always along a great circle.

Some examples of great circles on the celestial sphere include: the Horizon, the Celestial Equator, and the Ecliptic.

The Horizon is the line that separates Earth from Sky. More precisely, it is the line that divides all of the directions you can possibly look into two categories: those which intersect the Earth, and those which do not. At many locations, the Horizon is obscured by trees, buildings, mountains etc.. However, if you are on a ship at sea, the Horizon is strikingly apparent.

The horizon is the Fundamental Plane of the Horizontal Coordinate System. In other words, it is the locus of points which have an Altitude of zero degrees.

As explained in the Sidereal Time article, the Right Ascension of an object indicates the Sidereal Time at which it will transit across your Local Meridian. An object's Hour Angle is defined as the difference between the current Local Sidereal Time and the Right Ascension of the object:

HA_obj = LST - RA_obj

Thus, the object's Hour Angle indicates how much Sidereal Time has passed since the object was on the Local Meridian. It is also the angular distance between the object and the meridian, measured in hours (1 hour = 15 degrees). For example, if an object has an hour angle of 2.5 hours, it transited across the Local Meridian 2.5 hours ago, and is currently 37.5 degrees West of the Meridian. Negative Hour Angles indicate the time until the next transit across the Local Meridian. Of course, an Hour Angle of zero means the object is currently on the Local Meridian.

The Local Meridian is an imaginary Great Circle on the Celestial Sphere that is perpendicular to the local Horizon. It passes through the North point on the Horizon, through the Celestial Pole, up to the Zenith, and through the South point on the Horizon.

Because it is fixed to the local Horizon, stars will appear to drift past the Local Meridian as the Earth spins. You can use an object's Right Ascension and the Local Sidereal Time to determine when it will cross your Local Meridian (see Hour Angle).

Precession is the gradual change in the direction of the Earth's spin axis. The spin axis traces a cone, completing a full circuit in 26,000 years. If you have ever spun a top or a dreidel, the “wobbling” rotation of the top as it spins is precession.

Because the direction of the Earth's spin axis changes, so does the location of the Celestial Poles.

The reason for the Earth's precession is complicated. The Earth is not a perfect sphere, it is a bit flattened, meaning the Great Circle of the equator is longer than a “meridonal” great circle that passes through the poles. Also, the Moon and Sun lie outside the Earth's Equatorial plane. As a result, the gravitational pull of the Moon and Sun on the oblate Earth induces a slight torque in addition to a linear force. This torque on the spinning body of the Earth leads to the precessional motion.

The Zenith is the point in the sky where you are looking when you look “straight up” from the ground. More precisely, it is the point on the sky with an Altitude of +90 Degrees; it is the Pole of the Horizontal Coordinate System. Geometrically, it is the point on the Celestial Sphere intersected by a line drawn from the centre of the Earth through your location on the Earth's surface.

The Zenith is, by definition, a point along the Local Meridian.

Julian Days are a way of reckoning the current date by a simple count of the number of days that have passed since some remote, arbitrary date. This number of days is called the Julian Day, abbreviated as JD. The starting point, JD=0, is January 1, 4713 BC (or -4712 January 1, since there was no year '0'). Julian Days are very useful because they make it easy to determine the number of days between two events by simply subtracting their Julian Day numbers. Such a calculation is difficult for the standard (Gregorian) calendar, because days are grouped into months, which contain a variable number of days, and there is the added complication of Leap Years.

Converting from the standard (Gregorian) calendar to Julian Days and vice versa is best left to a special program written to do this, such as the KStars Astrocalculator. However, for those interested, here is a simple example of a Gregorian to Julian day converter:

JD = D - 32075 + 1461*( Y + 4800 + ( M - 14 ) / 12 ) / 4 + 367*( M - 2 - ( M - 14 ) / 12 * 12 ) / 12 - 3*( ( Y + 4900 + ( M - 14 ) / 12 ) / 100 ) / 4

where D is the day (1-31), M is the Month (1-12), and Y is the year (1801-2099). Note that this formula only works for dates between 1801 and 2099. More remote dates require a more complicated transformation.

An example Julian Day is: JD 2440588, which corresponds to 1 Jan, 1970.

Julian Days can also be used to tell time; the time of day is expressed as a fraction of a full day, with 12:00 noon (not midnight) as the zero point. So, 3:00 pm on 1 Jan 1970 is JD 2440588.125 (since 3:00 pm is 3 hours since noon, and 3/24 = 0.125 day). Note that the Julian Day is always determined from Universal Time, not Local Time.

Astronomers use certain Julian Day values as important reference points, called Epochs. One widely-used epoch is called J2000; it is the Julian Day for 1 Jan, 2000 at 12:00 noon = JD 2451545.0.

Much more information on Julian Days is available on the internet. A good starting point is the U.S. Naval Observatory. If that site is not available when you read this, try searching for “Julian Day” with your favourite search engine.

The Earth has two major components of motion. First, it spins on its rotational axis; a full spin rotation takes one Day to complete. Second, it orbits around the Sun; a full orbital rotation takes one Year to complete.

There are normally 365 days in one calendar year, but it turns out that a true year (i.e., a full orbit of the Earth around the Sun; also called a tropical year) is a little bit longer than 365 days. In other words, in the time it takes the Earth to complete one orbital circuit, it completes 365.24219 spin rotations. Do not be too surprised by this; there is no reason to expect the spin and orbital motions of the Earth to be synchronised in any way. However, it does make marking calendar time a bit awkward....

What would happen if we simply ignored the extra 0.24219 rotation at the end of the year, and simply defined a calendar year to always be 365.0 days long? The calendar is basically a charting of the Earth's progress around the Sun. If we ignore the extra bit at the end of each year, then with every passing year, the calendar date lags a little more behind the true position of Earth around the Sun. In just a few decades, the dates of the solstices and equinoxes will have drifted noticeably.

In fact, it used to be that all years were defined to have 365.0 days, and the calendar “drifted” away from the true seasons as a result. In the year 46 BCE, Julius Caeser established the Julian Calendar, which implemented the world's first leap years: He decreed that every 4th year would be 366 days long, so that a year was 365.25 days long, on average. This basically solved the calendar drift problem.

However, the problem wasn't completely solved by the Julian calendar, because a tropical year isn't 365.25 days long; it's 365.24219 days long. You still have a calendar drift problem, it just takes many centuries to become noticeable. And so, in 1582, Pope Gregory XIII instituted the Gregorian calendar, which was largely the same as the Julian Calendar, with one more trick added for leap years: even Century years (those ending with the digits “00”) are only leap years if they are divisible by 400. So, the years 1700, 1800, and 1900 were not leap years (though they would have been under the Julian Calendar), whereas the year 2000 was a leap year. This change makes the average length of a year 365.2425 days. So, there is still a tiny calendar drift, but it amounts to an error of only 3 days in 10,000 years. The Gregorian calendar is still used as a standard calendar throughout most of the world.

Sidereal Time literally means “star time”. The time we are used to using in our everyday lives is Solar Time. The fundamental unit of Solar Time is a Day: the time it takes the Sun to travel 360 degrees around the sky, due to the rotation of the Earth. Smaller units of Solar Time are just divisions of a Day:

However, there is a problem with Solar Time. The Earth does not actually spin around 360 degrees in one Solar Day. The Earth is in orbit around the Sun, and over the course of one day, it moves about one Degree along its orbit (360 degrees/365.25 Days for a full orbit = about one Degree per Day). So, in 24 hours, the direction toward the Sun changes by about a Degree. Therefore, the Earth has to spin 361 degrees to make the Sun look like it has travelled 360 degrees around the Sky.

In astronomy, we are concerned with how long it takes the Earth to spin with respect to the “fixed” stars, not the Sun. So, we would like a timescale that removes the complication of Earth's orbit around the Sun, and just focuses on how long it takes the Earth to spin 360 degrees with respect to the stars. This rotational period is called a Sidereal Day. On average, it is 4 minutes shorter than a Solar Day, because of the extra 1 degree the Earth spins in a Solar Day. Rather than defining a Sidereal Day to be 23 hours, 56 minutes, we define Sidereal Hours, Minutes and Seconds that are the same fraction of a Day as their Solar counterparts. Therefore, one Solar Second = 1.00278 Sidereal Seconds.

The Sidereal Time is useful for determining where the stars are at any given time. Sidereal Time divides one full spin of the Earth into 24 Sidereal Hours; similarly, the map of the sky is divided into 24 Hours of Right Ascension. This is no coincidence; Local Sidereal Time (LST) indicates the Right Ascension on the sky that is currently crossing the Local Meridian. So, if a star has a Right Ascension of 05h 32m 24s, it will be on your meridian at LST=05:32:24. More generally, the difference between an object's RA and the Local Sidereal Time tells you how far from the Meridian the object is. For example, the same object at LST=06:32:24 (one Sidereal Hour later), will be one Hour of Right Ascension west of your meridian, which is 15 degrees. This angular distance from the meridian is called the object's Hour Angle.

The Earth is round, and it is always half-illuminated by the Sun. However, because the Earth is spinning, the half that is illuminated is always changing. We experience this as the passing of days wherever we are on the Earth's surface. At any given instant, there are places on the Earth passing from the dark half into the illuminated half (which is seen as dawn on the surface). At the same instant, on the opposite side of the Earth, points are passing from the illuminated half into darkness (which is seen as dusk at those locations). So, at any given time, different places on Earth are experiencing different parts of the day. Thus, Solar time is defined locally, so that the clock time at any location describes the part of the day consistently.

This localisation of time is accomplished by dividing the globe into 24 vertical slices called Time Zones. The Local Time is the same within any given zone, but the time in each zone is one Hour earlier than the time in the neighboring Zone to the East. Actually, this is a idealised simplification; real Time Zone boundaries are not straight vertical lines, because they often follow national boundaries and other political considerations.

Note that because the Local Time always increases by an hour when moving between Zones to the East, by the time you move through all 24 Time Zones, you are a full day ahead of where you started. We deal with this paradox by defining the International Date Line, which is a Time Zone boundary in the Pacific Ocean, between Asia and North America. Points just to the East of this line are 24 hours behind the points just to the West of the line. This leads to some interesting phenomena. A direct flight from Australia to California arrives before it departs. Also, the islands of Fiji straddle the International Date Line, so if you have a bad day on the West side of Fiji, you can go over to the East side of Fiji and have a chance to live the same day all over again.

The time on our clocks is essentially a measurement of the current position of the Sun in the sky, which is different for places at different Longitudes because the Earth is round (see Time Zones).

However, it is sometimes necessary to define a global time, one that is the same for all places on Earth. One way to do this is to pick a place on the Earth, and adopt the Local Time at that place as the Universal Time, abbreviated UT. (The name is a bit of a misnomer, since Universal Time has little to do with the Universe. It would perhaps be better to think of it as global time).

The geographic location chosen to represent Universal Time is Greenwich, England. The choice is arbitrary and historical. Universal Time became an important concept when European ships began to sail the wide open seas, far from any known landmarks. A navigator could reckon the ship's longitude by comparing the Local Time (as measured from the Sun's position) to the time back at the home port (as kept by an accurate clock on board the ship). Greenwich was home to England's Royal Observatory, which was charged with keeping time very accurately, so that ships in port could re-calibrate their clocks before setting sail.

A blackbody refers to an opaque object that emits thermal radiation. A perfect blackbody is one that absorbs all incoming light and does not reflect any. At room temperature, such an object would appear to be perfectly black (hence the term blackbody). However, if heated to a high temperature, a blackbody will begin to glow with thermal radiation.

In fact, all objects emit thermal radiation (as long as their temperature is above Absolute Zero, or -273.15 degrees Celsius), but no object emits thermal radiation perfectly; rather, they are better at emitting/absorbing some wavelengths of light than others. These uneven efficiencies make it difficult to study the interaction of light, heat and matter using normal objects.

Fortunately, it is possible to construct a nearly-perfect blackbody. Construct a box made of a thermally conductive material, such as metal. The box should be completely closed on all sides, so that the inside forms a cavity that does not receive light from the surroundings. Then, make a small hole somewhere on the box. The light coming out of this hole will almost perfectly resemble the light from an ideal blackbody, for the temperature of the air inside the box.

At the beginning of the 20th century, scientists Lord Rayleigh, and Max Planck (among others) studied the blackbody radiation using such a device. After much work, Planck was able to empirically describe the intensity of light emitted by a blackbody as a function of wavelength. Furthermore, he was able to describe how this spectrum would change as the temperature changed. Planck's work on blackbody radiation is one of the areas of physics that led to the foundation of the wonderful science of Quantum Mechanics, but that is unfortunately beyond the scope of this article.

What Planck and the others found was that as the temperature of a blackbody increases, the total amount of light emitted per second increases, and the wavelength of the spectrum's peak shifts to bluer colours (see Figure 1).

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Figure 1

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For example, an iron bar becomes orange-red when heated to high temperatures and its colour progressively shifts toward blue and white as it is heated further.

In 1893, German physicist Wilhelm Wien quantified the relationship between blackbody temperature and the wavelength of the spectral peak with the following equation:

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where T is the temperature in Kelvin. Wien's law (also known as Wien's displacement law) states that the wavelength of maximum emission from a blackbody is inversely proportional to its temperature. This makes sense; shorter-wavelength (higher-frequency) light corresponds to higher-energy photons, which you would expect from a higher-temperature object.

For example, the sun has an average temperature of 5800 K, so its wavelength of maximum emission is given by:

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This wavelengths falls in the green region of the visible light spectrum, but the sun's continuum radiates photons both longer and shorter than lambda(max) and the human eyes perceives the sun's colour as yellow/white.

In 1879, Austrian physicist Stephan Josef Stefan showed that the luminosity, L, of a black body is proportional to the 4th power of its temperature T.

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where A is the surface area, alpha is a constant of proportionality, and T is the temperature in Kelvin. That is, if we double the temperature (e.g. 1000 K to 2000 K) then the total energy radiated from a blackbody increase by a factor of 2^4 or 16.

Five years later, Austrian physicist Ludwig Boltzman derived the same equation and is now known as the Stefan-Boltzman law. If we assume a spherical star with radius R, then the luminosity of such a star is

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where R is the star radius in cm, and the alpha is the Stefan-Boltzman constant, which has the value:

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Scientists are now quite comfortable with the idea that 90% of the mass is the universe is in a form of matter that cannot be seen.

Despite comprehensive maps of the nearby universe that cover the spectrum from radio to gamma rays, we are only able to account of 10% of the mass that must be out there. As Bruce H. Margon, an astronomer at the University of Washington, told the New York Times in 2001: [It's a fairly embarrassing situation to admit that we can't find 90 percent of the universe].

The term given this “missing mass” is Dark Matter, and those two words pretty well sum up everything we know about it at this point. We know there is “Matter”, because we can see the effects of its gravitational influence. However, the matter emits no detectable electromagnetic radiation at all, hence it is “Dark”. There exist several theories to account for the missing mass ranging from exotic subatomic particles, to a population of isolated black holes, to less exotic brown and white dwarfs. The term “missing mass” might be misleading, since the mass itself is not missing, just its light. But what is exactly dark matter and how do we really know it exists, if we cannot see it?

The story began in 1933 when Astronomer Fritz Zwicky was studying the motions of distant and massive clusters of galaxies, specifically the Coma cluster and the Virgo cluster. Zwicky estimated the mass of each galaxy in the cluster based on their luminosity, and added up all of the galaxy masses to get a total cluster mass. He then made a second, independent estimate of the cluster mass, based on measuring the spread in velocities of the individual galaxies in the cluster. To his suprise, this second dynamical mass estimate was 400 times larger than the estimate based on the galaxy light.

Although the evidence was strong at Zwicky's time, it was not until the 1970s that scientists began to explore this discrepancy comprehensively. It was at this time that the existence of Dark Matter began to be taken seriously. The existence of such matter would not only resolve the mass deficit in galaxy clusters; it would also have far more reaching consequences for the evolution and fate of the universe itself.

Another phenomenon that suggested the need for dark matter is the rotational curves of Spiral Galaxies. Spiral Galaxies contain a large population of stars that orbit the Galactic centre on nearly circular orbits, much like planets orbit a star. Like planetary orbits, stars with larger galactic orbits are expected to have slower orbital speeds (this is just a statement of Kepler's 3rd Law). Actually, Kepler's 3rd Law only applies to stars near the perimeter of a Spiral Galaxy, because it assumes the mass enclosed by the orbit to be constant.

However, astronomers have made observations of the orbital speeds of stars in the outer parts of a large number of spiral galaxies, and none of them follow Kepler's 3rd Law as expected. Instead of falling off at larger radii, the orbital speeds remain remarkably constant. The implication is that the mass enclosed by larger-radius orbits increases, even for stars that are apparently near the edge of the galaxy. While they are near the edge of the luminous part of the galaxy, the galaxy has a mass profile that apparently continues well beyond the regions occupied by stars.

Here is another way to think about it: Consider the stars near the perimeter of a spiral galaxy, with typical observed orbital velocities of 200 kilometers per second. If the galaxy consisted of only the matter that we can see, these stars would very quickly fly off from the galaxy, because their orbital speeds are four times larger than the galaxy's escape velocity. Since galaxies are not seen to be spinning apart, there must be mass in the galaxy that we are not accounting for when we add up all the parts we can see.

Several theories have surfaced in literature to account for the missing mass such as WIMPs (Weakly Interacting Massive Particles), MACHOs (MAssive Compact Halo Objects), primordial black holes, massive neutrinos, and others; each with their pros and cons. No single theory has yet been accepted by the astronomical community, because we so far lack the means to conclusively test one theory against the other.

The flux is the amount of energy that passes through a unit area each second.

Astronomers use flux to denote the apparent brightness of a celestial body. The apparent brightness is defined as the the amount of light received from a star above the earth atmosphere passing through a unit area each second. Therefore, the apparent brightness is simply the flux we receive from a star.

The flux measures the rate of flow of energy that passes through each cm^2 (or any unit area) of an object's surface each second. The detected flux depends on the distance from the source that radiates the energy. This is because the energy has to spread over a volume of space before it reaches us. Let us assume that we have an imaginary balloon that encloses a star. Each dot on the balloon represents a unit of energy emitted from the star. Initially, the dots in an area of one cm^2 are in close proximity to each other and the flux (energy emitted per square centimetre per second) is high. After a distance d, the volume and surface area of the balloon increased causing the dots to spread away from each. Consequently, the number of dots (or energy) enclosed in one cm^2 has decreased as illustrated in Figure 1.

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Figure 1

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The flux is inversely proportional to distance by a simple r^2 relation. Therefore, if the distance is doubled, we receive 1/2^2 or 1/4th of the original flux. From a fundamental standpoint, the flux is the luminosity per unit area:

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where (4 * PI * R^2) is the surface area of a sphere (or a balloon!) with a radius R. Flux is measured in Watts/m^2/s or as commonly used by astronomers: Ergs/cm^2/s. For example, the luminosity of the sun is L = 3.90 * 10^26 W. That is, in one second the sun radiates 3.90 * 10^26 joules of energy into space. Thus, the flux we receive passing through one square centimetre from the sun at a distance of one AU (1.496 * 10^13 cm) is:

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Luminosity is the amount of energy emitted by a star each second.

All stars radiate light over a broad range of frequencies in the electromagnetic spectrum from the low energy radio waves up to the highly energetic gamma rays. A star that emits predominately in the ultra-violet region of the spectrum produces a total amount of energy magnitudes larger than that produced in a star that emits principally in the infrared. Therefore, luminosity is a measure of energy emitted by a star over all wavelengths. The relationship between wavelength and energy was quantified by Einstein as E = h * v where v is the frequency, h is the Planck constant, and E is the photon energy in joules. That is, shorter wavelengths (and thus higher frequencies) correspond to higher energies.

For example, a wavelength of lambda = 10 meter lies in the radio region of the electromagnetic spectrum and has a frequency of f = c / lambda = 3 * 10^8 m/s / 10 = 30 MHz where c is the speed of light. The energy of this photon is E = h * v = 6.625 * 10^-34 J s * 30 Mhz = 1.988 * 10^-26 joules. On the other hand, visible light has much shorter wavelengths and higher frequencies. A photon that has a wavelength of lambda = 5 * 10^-9 meters (A greenish photon) has an energy E = 3.975 * 10^-17 joules which is over a billion times higher than the energy of a radio photon. Similarly, a photon of red light (wavelength lambda = 700 nm) has less energy than a photon of violet light (wavelength lambda = 400 nm).

Luminosity depends both on temperature and surface area. This makes sense because a burning log radiates more energy than a match, even though both have the same temperature. Similarly, an iron rod heated to 2000 degrees emits more energy than when it is heated to only 200 degrees.

Luminosity is a very fundamental quantity in Astronomy and Astrophysics. Much of what is learnt about celestial objects comes from analysing their light. This is because the physical processes that occur inside stars gets recorded and transmitted by light. Luminosity is measured in units of energy per second. Astronomers prefer to use Ergs, rather than Watts, when quantifying luminosity.

Parallax is the apparent change of an observed object's position caused by a shift in the observer's position. As an example, hold your hand in front of you at arm's length, and observe an object on the other side of the room behind your hand. Now tilt your head to your right shoulder, and your hand will appear on the left side of the distant object. Tilt your head to your left shoulder, and your hand will appear to shift to the right side of the distant object.

Because the Earth is in orbit around the Sun, we observe the sky from a constantly moving position in space. Therefore, we should expect to see an annual parallax effect, in which the positions of nearby objects appear to “wobble” back and forth in response to our motion around the Sun. This does in fact happen, but the distances to even the nearest stars are so great that you need to make careful observations with a telescope to detect it^[2].

Modern telescopes allow astronomers to use the annual parallax to measure the distance to nearby stars, using triangulation. The astronomer carefully measures the position of the star on two dates, spaced six months apart. The nearer the star is to the Sun, the larger the apparent shift in its position will be between the two dates.

Over the six-month period, the Earth has moved through half its orbit around the Sun; in this time its position has changed by 2 Astronomical Units (abbreviated AU; 1 AU is the distance from the Earth to the Sun, or about 150 million kilometers). This sounds like a really long distance, but even the nearest star to the Sun (alpha Centauri) is about 40 trillion kilometers away. Therefore, the annual parallax is very small, typically smaller than one arcsecond, which is only 1/3600 of one degree. A convenient distance unit for nearby stars is the parsec, which is short for "parallax arcsecond". One parsec is the distance a star would have if its observed parallax angle was one arcsecond. It is equal to 3.26 light-years, or 31 trillion kilometers^[3].

Retrograde Motion is the orbital motion of a body in a direction opposite that which is normal to spatial bodies within a given system.

When we observe the sky, we expect most objects to appear to move in a particular direction with the passing of time. The apparent motion of most bodies in the sky is from east to west. However it is possible to observe a body moving west to east, such as an artificial satellite or space shuttle that is orbiting eastward. This orbit is considered Retrograde Motion.

Retrograde Motion is most often used in reference to the motion of the outer planets (Mars, Jupiter, Saturn, and so forth). Though these planets appear to move from east to west on a nightly basis in response to the spin of the Earth, they are actually drifting slowly eastward with respect to the stationary stars, which can be observed by noting the position of these planets for several nights in a row. This motion is normal for these planets, however, and not considered Retrograde Motion. However, since the Earth completes its orbit in a shorter period of time than these outer planets, we occasionally overtake an outer planet, like a faster car on a multiple-lane highway. When this occurs, the planet we are passing will first appear to stop its eastward drift, and it will then appear to drift back toward the west. This is Retrograde Motion, since it is in a direction opposite that which is typical for planets. Finally, as the Earth swings past the the planet in its orbit, they appear to resume their normal west-to-east drift on successive nights.

This Retrograde Motion of the planets puzzled ancient Greek astronomers, and was one reason why they named these bodies “planets” which in Greek means “wanderers”.

Elliptical galaxies are spheroidal concentrations of billions of stars that resemble Globular Clusters on a grand scale. They have very little internal structure; the density of stars declines smoothly from the concentrated centre to the diffuse edge, and they can have a broad range of ellipticities (or aspect ratios). They typically contain very little interstellar gas and dust, and no young stellar populations (although there are exceptions to these rules). Edwin Hubble referred to Elliptical galaxies as “early-type” galaxies, because he thought that they evolved to become Spiral Galaxies (which he called “late-type” galaxies). Astronomers actually now believe the opposite is the case (i.e., that Spiral galaxies can turn into Elliptical galaxies), but Hubble's early- and late-type labels are still used.

Once thought to be a simple galaxy type, ellipticals are now known to be quite complex objects. Part of this complexity is due to their amazing history: ellipticals are thought to be the end product of the merger of two Spiral galaxies. You can view a computer simulation MPEG movie of such a merger at this NASA HST webpage (warning: the file is 3.4 MB).

Elliptical galaxies span a very wide range of sizes and luminosities, from giant Ellipticals hundreds of thousands of light years across and nearly a trillion times brighter than the sun, to dwarf Ellipticals just a bit brighter than the average globular cluster. They are divided to several morphological classes:

Spiral galaxies are huge collections of billions of stars, most of which are flattened into a disk shape, with a bright, spherical bulge of stars at its centre. Within the disk, there are typically bright arms where the youngest, brightest stars are found. These arms wind out from the centre in a spiral pattern, giving the galaxies their name. Spiral galaxies look a bit like hurricanes, or like water flowing down a drain. They are some of the most beautiful objects in the sky.

Galaxies are classified using a “tuning fork diagram”. The end of the fork classifies elliptical galaxies on a scale from the roundest, which is an E0, to those that appear most flattened, which is rated as E7. The “tines” of the tuning fork are where the two types of spiral galaxies are classified: normal spirals, and “barred” spirals. A barred spiral is one whose nuclear bulge is stretched out into a line, so it literally looks like it has a “bar” of stars in its centre.

Both types of spiral galaxies are sub-classified according to the prominence of their central “bulge” of stars, their overall surface brightness, and how tightly their spiral arms are wound. These characteristics are related, so that an Sa galaxy has a large central bulge, a high surface brightness, and tightly-wound spiral arms. An Sb galaxy has a smaller bulge, a dimmer disk, and looser arms than an Sa, and so on through Sc and Sd. Barred galaxies use the same classification scheme, indicated by types SBa, SBb, SBc, and SBd.

There is another class of galaxy called S0, which is morphologically a transitional type between true spirals and ellipticals. Its spiral arms are so tightly wound as to be indistinguishable; S0 galaxies have disks with a uniform brightness. They also have an extremely dominant bulge.

The Milky Way galaxy, which is home to earth and all of the stars in our sky, is a Spiral Galaxy, and is believed to be a barred spiral. The name “Milky Way” refers to a band of very faint stars in the sky. This band is the result of looking in the plane of our galaxy's disk from our perspective inside it.

Spiral galaxies are very dynamic entities. They are hotbeds of star formation, and contain many young stars in their disks. Their central bulges tend to be made of older stars, and their diffuse halos are made of the very oldest stars in the Universe. Star formation is active in the disks because that is where the gas and dust are most concentrated; gas and dust are the building blocks of star formation.

Modern telescopes have revealed that many Spiral galaxies harbour supermassive black holes at their centres, with masses that can exceed that of a billion Suns. Both elliptical and spiral galaxies are known to contain these exotic objects; in fact many astronomers now believe that all large galaxies contain a supermassive black hole in their nucleus. Our own Milky Way is known to harbour a black hole in its core with a mass millions of times bigger than a star's mass.

2500 years ago, the ancient Greek astronomer Hipparchus classified the brightnesses of visible stars in the sky on a scale from 1 to 6. He called the very brightest stars in the sky “first magnitude”, and the very faintest stars he could see “sixth magnitude”. Amazingly, two and a half millenia later, Hipparchus's classification scheme is still widely used by astronomers, although it has since been modernised and quantified.

The modern magnitude scale is a quantitative measurement of the flux of light coming from a star, with a logarithmic scaling:

m = m_0 - 2.5 log (F / F_0)

If you do not understand the maths, this just says that the magnitude of a given star (m) is different from that of some standard star (m_0) by 2.5 times the logarithm of their flux ratio. The 2.5 *log factor means that if the flux ratio is 100, the difference in magnitudes is 5 mag. So, a 6th magnitude star is 100 times fainter than a 1st magnitude star. The reason Hipparchus's simple classification translates to a relatively complex function is that the human eye responds logarithmically to light.

There are several different magnitude scales in use, each of which serves a different purpose. The most common is the apparent magnitude scale; this is just the measure of how bright stars (and other objects) look to the human eye. The apparent magnitude scale defines the star Vega to have magnitude 0.0, and assigns magnitudes to all other objects using the above equation, and a measure of the flux ratio of each object to Vega.

It is difficult to understand stars using just the apparent magnitudes. Imagine two stars in the sky with the same apparent magnitude, so they appear to be equally bright. You cannot know just by looking if the two have the same intrinsic brightness; it is possible that one star is intrinsically brighter, but further away. If we knew the distances to the stars (see the parallax article), we could account for their distances and assign Absolute magnitudes which would reflect their true, intrinsic brightness. The absolute magnitude is defined as the apparent magnitude the star would have if observed from a distance of 10 parsecs (1 parsec is 3.26 light-years, or 3.1 x 10^18 cm). The absolute magnitude (M) can be determined from the apparent magnitude (m) and the distance in parsecs (d) using the formula:

M = m + 5 - 5 * log(d) (note that M=m when d=10).

The modern magnitude scale is no longer based on the human eye; it is based on photographic plates and photoelectric photometers. With telescopes, we can see objects much fainter than Hipparchus could see with his unaided eyes, so the magnitude scale has been extended beyond 6th magnitude. In fact, the Hubble Space Telescope can image stars nearly as faint as 30th magnitude, which is one trillion times fainter than Vega.

A final note: the magnitude is usually measured through a colour filter of some kind, and these magnitudes are denoted by a subscript describing the filter (i.e., m_V is the magnitude through a “visual” filter, which is greenish; m_B is the magnitude through a blue filter; m_pg is the photographic plate magnitude etc.).

Stars appear to be exclusively white at first glance. But if we look carefully, we can notice a range of colours: blue, white, red and even gold. In the winter constellation of Orion, a beautiful contrast is seen between the red Betelgeuse at Orion's "armpit" and the blue Bellatrix at the shoulder. What causes stars to exhibit different colours remained a mystery until two centuries ago, when Physicists gained enough understanding of the nature of light and the properties of matter at immensely high temperatures.

Specifically, it was the physics of blackbody radiation that enabled us to understand the variation of stellar colours. Shortly after blackbody radiation was understood, it was noticed that the spectra of stars look extremely similar to blackbody radiation curves of various temperatures, ranging from a few thousand Kelvin to ~50,000 Kelvin. The obvious conclusion is that stars are similar to blackbodies, and that the colour variation of stars is a direct consequence of their surface temperatures.

Cool stars (i.e., Spectral Type K and M) radiate most of their energy in the red and infrared region of the electromagnetic spectrum and thus appear red, while hot stars (i.e., Spectral Type O and B) emit mostly at blue and ultra-violet wavelengths, making them appear blue or white.

To estimate the surface temperature of a star, we can use the known relationship between the temperature of a blackbody, and the wavelength of light where its spectrum peaks. That is, as you increase the temperature of a blackbody, the peak of its spectrum moves to shorter (bluer) wavelengths of light. This is illustrated in Figure 1 where the intensity of three hypothetical stars is plotted against wavelength. The "rainbow" indicates the range of wavelengths that are visible to the human eye.

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Figure 1

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This simple method is conceptually correct, but it cannot be used to obtain stellar temperatures accurately, because stars are not perfect blackbodies. The presence of various elements in the star's atmosphere will cause certain wavelengths of light to be absorbed. Because these absorption lines are not uniformly distributed over the spectrum, they can skew the position of the spectral peak. Moreover, obtaining a usable spectrum of a star is a time-intensive process and is prohibitively inefficient for large samples of stars.

An alternative method utilises photometry to measure the intensity of light passing through different filters. Each filter allows only a specific part of the spectrum of light to pass through while rejecting all others. A widely used photometric system is called the Johnson UBV system. It employs three bandpass filters: U ("Ultra-violet"), B ("Blue"), and V ("Visible"); each occupying different regions of the electromagnetic spectrum.

The process of UBV photometry involves using light sensitive devices (such as film or CCD cameras) and aiming a telescope at a star to measure the intensity of light that passes through each of the filters individually. This procedure gives three apparent brightnesses or fluxes (amount of energy per cm^2 per second) designated by Fu, Fb, and Fv. The ratio of fluxes Fu/Fb and Fb/Fv is a quantitative measure of the star's "colour", and these ratios can be used to establish a temperature scale for stars. Generally speaking, the larger the Fu/Fb and Fb/Fv ratios of a star, the hotter its surface temperature.

For example, the star Bellatrix in Orion has Fb/Fv = 1.22, indicating that it is brighter through the B filter than through the V filter. furthermore, its Fu/Fb ratio is 2.22, so it is brightest through the U filter. This indicates that the star must be very hot indeed, since the position of its spectral peak must be somewhere in the range of the U filter, or at an even shorter wavelength. The surface temperature of Bellatrix (as determined from comparing its spectrum to detailed models that account for its absorption lines) is about 25,000 Kelvin.

We can repeat this analysis for the star Betelgeuse. Its Fb/Fv and Fu/Fb ratios are 0.15 and 0.18, respectively, so it is brightest in V and dimmest in U. So, the spectral peak of Betelgeuse must be somewhere in the range of the V filter, or at an even longer wavelength. The surface temperature of Betelgeuse is only 2,400 Kelvin.

Astronomers prefer to express star colours in terms of a difference in magnitudes, rather than a ratio of fluxes. Therefore, going back to blue Bellatrix we have a colour index equal to

B - V = -2.5 log (Fb/Fv) = -2.5 log (1.22) = -0.22,

Similarly, the colour index for red Betelgeuse is

B - V = -2.5 log (Fb/Fv) = -2.5 log (0.18) = 1.85

The colour indices, like the magnitude scale, run backward. Hot and blue stars have smaller and negative values of B-V than the cooler and redder stars.

An Astronomer can then use the colour indices for a star, after correcting for reddening and interstellar extinction, to obtain an accurate temperature of that star. The relationship between B-V and temperature is illustrated in Figure 2.

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Figure 2

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The Sun with surface temperature of 5,800 K has a B-V index of 0.62.

The KStars Astrocalculator provides several modules that give you direct access to algorithms used by the program. The modules are organised by subject:

Angular Distance module

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The Angular Distance tool is used to measure the angle between any two points on the sky. You simply specify the Equatorial coordinates of the desired pair of points, and then press the Compute button to obtain the angle between the two points.

There is also a Batch mode for this module. In batch mode, you specify an input filename which contains four numbers per line: the RA and Dec values for pairs of points. Alternatively, you can specify a single value for any of these four coordinates in the calculator panel (the corresponding values in the input file should be skipped if they are specified in the calculator).

Once you have specified the input filename and an output filename, simply press the Run button to generate the output file.

Apparent Coordinates module

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The Apparent Coordinates module converts the catalogue coordinates of a point in the sky to its apparent coordinates for any date. The coordinates of objects in the sky are not fixed, because of precession, nutation and aberration. This module takes these effects into account.

To use the module, first enter the desired target date and time in the Target Time/Date section. Then, enter the catalogue coordinates in the Catalog Coordinates section. You can also specify the catalogue's epoch here (usually 2000.0 for modern object catalogues). Finally, press the Compute button, and the object's coordinates for the target date will be displayed in the Apparent Coordinates section.

Ecliptic Coordinates module

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This module converts between Equatorial coordinates and Ecliptic coordinates. First, select which coordinates should be taken as input values in the Choose Input Coordinates section. Then, fill in the corresponding coordinate values in either the Ecliptic coordinates or Equatorial coordinates section. Finally, press the Compute button, and the complementary coordinates will be filled in.

The module contains a batch mode for converting several coordinate pairs at once. You must construct an input file in which each line contains two values: the input coordinate pairs (either Equatorial or Ecliptic). Then specify which coordinates you are using as input, and identify the input and output filenames. Finally, press the Run button to generate the output file, which will contain the converted coordinates (Equatorial or Ecliptic; the complement of what you chose as the input values).

Equatorial/Galactic Coordinates module

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This module converts from Equatorial coordinates to Galactic coordinates, and vice versa. First, select which coordinates should be taken as input values in the Input Selection section. Then, fill in the corresponding coordinate values in either the Galactic coordinates or Equatorial coordinates section. Finally, press the Compute button, and the complementary coordinates will be filled in.

Horizontal Coordinates module

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This module converts from Equatorial coordinates to Horizontal coordinates. First, select the date, time, and geographic coordinates for the calculation in the Input Data section. Then, fill in the equatorial coordinates to be converted and their catalogue epoch in the Equatorial Coordinates section. When you press the Compute button, the corresponding Horizontal coordinates will be presented in the Horizontal Coordinates section.

Precession module

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This module is similar to the Apparent Coordinates module, but it only applies the effect of precession, not of nutation or aberration.

To use the module, first enter the input coordinates and their epoch in the Original Coordinates section. You must also fill in the target epoch in the Precessed Coordinates section. Then, press the Compute button, and the object's coordinates, precessed to the target Epoch, are presented in the Precessed Coordinates section.

Geodetic Coordinates module

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The normal geographic coordinate system assumes that the Earth is a perfect sphere. This is nearly true, so for most purposes geographic coordinates are fine. If very high precision is required, then we must take the true shape of the Earth into account. The Earth is an ellipsoid; the distance around the equator is about 0.3% longer than a Great Circle that passes through the poles. The Geodetic Coordinate system takes this ellipsoidal shape into account, and expresses the position on the Earth's surface in Cartesian coordinates (X, Y, and Z).

To use the module, first select which coordinates you will use as input in the Input Selection section. Then, fill in the input coordinates in either the Cartesian Coordinates section or the Geographic Coordinates section. When you press the Compute button, the corresponding coordinates will be filled in.

Planet Coordinates module

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The Planet Coordinates module computes positional data for any major solar system body, for any time and date and any geographic location. Simply select the solar system body from the drop-down list, and specify the desired date, time and geographic coordinates (these values are preset to the current KStars settings). Then press the Compute button to determine the Equatorial, Horizontal and Ecliptic coordinates of the body.

There is a batch mode for this module. You must construct an input file in which each line specifies values for the input parameters (solar system body, date, time, longitude, and latitude). You may choose to specify a constant value for some of the parameters in the calculator window (these parameters should be skipped in the input file). You may also specify which of the output parameters (Equatorial, Horizontal and Ecliptic coordinates) should be calculated. Finally, specify the input and output filenames, and press the Run button to generate the output file with the computed values.

Day Duration module

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This module computes the length of day as well as sunrise, sun-transit (noon), and sunset times for any calendar date, for any location on Earth. First fill in the desired geographic coordinates and date, then press the Compute button.

Equinoxes and Solstices module

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The Equinoxes and Solstices module calculates the date and time of an equinox or solstice for a given year. You specify which event (Spring Equinox, Summer Solstice, Autumn Equinox or Winter Solstice) should be investigated, and the year. Then press the Compute button to obtain the date and time of the event, and the length of the corresponding season, in days.

There is a batch mode for this module. To use it, simply generate an input file whose lines each contain a year for which the Equinox and Solstice data will be computed. Then specify the input and output filenames, and press the Run button to generate the output file. Each line in the output file contains the input year, the date and time of each event, and the length of each season.

Julian Day module

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This module converts between the calendar date, the Julian Day and the Modified Julian Day. The Modified Julian Day is simply equal to the Julian Day - 2,400,000.5.

To use the module, select which of the three dates will be the input, and then fill in its value. Then press the Compute button, and the corresponding values for the other two date systems will be displayed.

Tip

Exercise:

What calendar date does MJD = 0.0 correspond to?

Sidereal Time module

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This module converts between Universal Time and Local Sidereal Time. First, select whether you will use Universal Time or Sidereal Time as an input value in the Input Selection section. You must also specify a geographic longitude, and a date for the calculation, in addition to either the Universal Time or the Sidereal Time value. When you press the Compute button, the corresponding value for the other Time will be displayed.

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This tool plots the altitude of any objects as a function of time, for any date and location on Earth. The top section is a graphical plot of altitude angle on the vertical axis, and time on the horizontal axis. The time is shown both as standard local time along the bottom, and sidereal time along the top. The bottom half of the graph is shaded green to indicate that points in this region are below the horizon.

There are a few ways to add curves to the plot. The simplest way to add the curve of an existing object is to simply type its name in the Name input field, and press Enter, or the Plot button. If the text you enter is found in the object database, the object's curve is added to the graph. You can also press the Browse button to open the Find Object Window to select an object from the list of known objects. If you want to add a point that does not exist in the object database, simply enter a name for the point, and then fill in the coordinates in the RA and Dec input fields. Then press the Plot button to add the curve for your custom object to the plot (note that you have to pick a name that does not already exist in the object database for this to work).

When you add an object to the plot, its altitude vs. time curve is plotted with a thick white line, and its name is added to the listbox at the lower right. Any objects that were already present are plotted with a thinner red curve. You can choose which object is plotted with the thick white curve by highlighting its name in the listbox.

These curves show the objects' Altitude (angle above the horizon) as a function of time. When a curve passes from the lower half to the upper half, the object has risen; when it falls back to the lower half, it has set. For example, in the screenshot, the minor planet Quaoar is rising at around 15:30 local time, and is setting at about 00:30.

The Altitude of an object depends on both where you are on Earth, and on the Date. By default, the Tool adopts the Location and Date from the current KStars settings. You can change these parameters in the Date & Location Tab. To change the Location, you can press the Choose City... button to open the Set Geographic Location Window, or enter Longitude and Latitude values manually in the input fields, and press the Update button. To change the Date, use the Date picker widget, then press Update. Note that any curves you had already plotted will be automatically updated when you change the Date and/or Location.

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The “What's Up Tonight?” (WUT) tool displays a list of objects that will be visible at night from any location, on any date. By default, the Date and Location are taken from the current settings in the main window, but you can change either value using the Change Date and Change Location buttons at the top of the WUT window.

The WUT tool also displays a short almanac of data for the selected date: the rise and set times for the Sun and moon, the duration of the night and the Moon's illumination fraction.

Below the almanac is where the object information is displayed. The objects are organised into type categories. Select an object type in the box labelled Choose a Category, and all objects of that type which are above the horizon on the selected night will be displayed in the box labelled Matching Objects. For example, in the screenshot, the Planets category has been selected, and four planets which are up on the selected night are displayed (Mars, Neptune, Pluto, and Uranus). When an object in the list is selected, its rise, set and transit times are displayed in the lower-right panel. In addition, you can press the Object Details... button to open the Object Details window for that object.

By default, the WUT will display objects which are above the horizon between sunset and midnight (i.e., “in the evening”). You can choose to show objects which are up between midnight and dawn (“in the morning”), or between dusk and dawn (“any time tonight”) using the combobox near the top of the window.

KDE applications can be controlled externally from another program, from a console prompt, or from a shell script using the Desktop COmmunication Protocol (DCOP). KStars takes advantage of this feature to allow rather complex behaviours to be scripted and played back at any time. This can be used, for example, to create a classroom demo to illustrate an astronomical concept.

The problem with DCOP scripts is, writing them is a bit like programming, and can seem a daunting task to those who do not have programming experience. The Script Builder Tool provides a GUI point-and-click interface for constructing KStars DCOP scripts, making it very easy to create complex scripts.

Introduction to the Script Builder

Before explaining how to use the Script Builder, I provide a very brief introduction to all of the GUI components; for more infomation, use the "What's This?" function.

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The Script Builder is shown in the above screenshot. The box on the left is the Current Script box; it shows the list of commands that comprise the current working script. The box on the right is the Function Browser; it displays the list of all available script functions. Below the Function Browser, there is a small panel which will display short documentation about the script function highlighted in the Function Browser. The panel below the Current Script box is the Function Arguments panel; when a function is highlighted in the Current Script box, this panel will contain items for specifying values for any arguments that the highlighted function requires.

Along the top of the window, there is a row of buttons which operate on the script as a whole. From left to right, they are: New Script, Open Script, Save Script, Save Script As..., and Test Script. The function of these buttons should be obvious, except perhaps the last button. Pressing Test Script will attempt to run the current script in the main KStars window. You should move the Script Builder window out of the way before pressing this, so you can see the results.

In the centre of the window, there is a column of buttons which operate on individual script functions. From top to bottom, they are: Add Function, Remove Function, Copy Function, Move Up, and Move Down. Add Function adds the currently-highlighted function in the Function Browser to the Current Script box (you can also add a function by double-clicking on it). The rest of the buttons operate on the function highlighted in the Current Script box, either removing it, duplicating it, or changing its position in the current script.

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This tool displays a model of our solar system as seen from above, for the current date and time in the main window. The Sun is drawn as a yellow dot in the centre of the plot, and the orbits of the planets are drawn as circles with correct relative diameters, centred on the Sun. The current position of each planet along its orbit is drawn as a coloured dot along with a name label. The display can be zoomed in and out with the + and - keys.

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This tool displays the positions of Jupiter's four largest moons (Io, Europa, Ganymede, and Callisto) relative to Jupiter, as a function of time. Time is plotted vertically; the units are days and “time=0.0” corresponds to the current simulation time. The horizontal axis displays the angular offset from Jupiter's position, in arcminutes. The offset is measured along the direction of Jupiter's equator. Each moon's position as a function of time traces a sinusoidal path in the plot, as the moon orbits around Jupiter. Each track is assigned a different colour to distinguish it from the others; the name labels at the top of the window indicate the colour used by each moon.

The plot can be manipulated with the keyboard. The time axis can be expanded or compressed using the + and - keys. The time displayed at the center of the window can be changed with the [ and ] keys.

Most telescopes are equipped with RS232 interface for remote control. Connect the RS232 jack in your telescope to your computer's Serial/USB port. Traditionally, the RS232 connects to the serial port of your computer, but since many new laptops abandoned the serial port in favour of USB/FireWire ports, you might need to obtain a Serial to USB adaptor to use with new laptops.

After connecting your telescope to the Serial/USB port, turn your telescope on. It is highly recommended that you download and install the latest firmware for your telescope controller.

The telescope needs to be aligned before it can be used properly. Align your telescope (one or two stars alignment) as illustrated in your telescope manual.

KStars needs to verify time and location settings before connecting to the telescope. This insures proper tracking and synchronisation between the telescope and KStars. The following steps will enable you to connect to a device that is connected to your computer. To connect and control remote devices, please refer to remote device control section.

You can use the Telescope Setup Wizard and it will verify all the required information in the process. It can automatically scan ports for attached telescopes. You can run the wizard by selecting Telescope Setup Wizard from the Devices menu.

Alternatively, you can connect to a local telescope by performing the following steps:

You are now ready to use the device features, KStars conveniently provides two interchangeable GUI interfaces for controlling telescopes:

You are not restricted on using one interface over another as they can be both used simultaneously. Actions from the Sky map are automatically reflected in the INDI Control Panel and vice versa.

To connect to your telescope, you can either select Connect from your device popup menu or alternatively, you can press Connect under your device tab in the INDI Control Panel.

KStars automatically updates the telescope's longitude, latitude, and time based on current settings in KStars. You can enable/disable these updates from Configure INDI dialogue under the Devices menu.

If KStars communicates successfully with the telescope, it will retrieve the current RA and DEC from the telescope and will display a crosshair on the sky map indicating the telescope position.

That is it: your telescope is ready to explore the heavens.

KStars supports Finger Lakes instruments CCDs and any Video4Linux compatible device. Philips webcam extended features are supported as well.

You can run CCD and Video Capture devices from the Device Manager in the Devices menu. Like all INDI devices, some of the device controls will be accessible from the skymap. The device can be controlled fully from the INDI Control Panel.

The standard format for image capture is FITS. Once an image is captured and downloaded, it will be displayed in the KStars FITSViewer. To capture a sequence of images, use the Capture Image Sequence tool from the Devices menu. This tool is inactive until you establish a connection to an image device.

The INDI control panel offers many device properties not accessible from the sky map. The properties offered differ from one device to another. Nevertheless, all properties share common features that constrains how they are displayed and used:

KStars provides a simple yet powerful layer for remote device control. A detailed description of the layer is described in the INDI white paper.

You need to configure both the server and client machines for remote control:

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After you add a host, right click on the host to Connect or Disconnect. If a connection is established, you can control the telescope from the Sky map or INDI Control Panel exactly as described in the local/server section. It is as easy at that.