Kuiper Belt Objects Advanced Lab

Transcription

1 Kuiper Belt Objects Advanced Lab Introduction To date, there are over 763 known Kuiper Belt Objects (KBOs), including Pluto and Charon, however, because they are so distant, faint, and relatively new to astronomy, little is known about these objects. In this lab, you will observe several KBOs to measure the colors of the KBO, and if possible, determine the rotation period from a light curve. In this lab you will make the measurements of several KBOs, over several nights. The main instrument you will be using for this experiment will be the Braeside Observatory. The Braeside telescope is a 16 inch telescope located in Flagstaff, AZ. A CCD, or Charged Coupled Device, situated at the end of the telescope converts the light which enters the telescope into the digital image. The CCD camera used on the Braeside telescope is called a SITe CCD. Each pixel making up the CCD chip is a light sensitive device which measures the amount of light falling on it. The plate scale for the camera is arcsecond/pixel which gives a field of view of 6 arcmin square. The CCD camera has several filters. Each filter selects out a region of light known as a passband. The filters are used to measure the color of astronomical objects. If an object emits more light in the blue than in green (known as visual light) then the object will appear brighter in the B filter than in the V filter. The filter wheel has 8 slots open for filters which are described in Table 1. The scientific merits of this lab are to (i) obtain photometric measurements of as many KBOs as possible and (ii) measure any variation in color/brightness. While a small portion of the known KBOs have been photometrically measured, there is still a debate whether or not there are two color classes. Periodic variations in brightness can be related to rotation period while variations in color would reveal hints to surface composition. In addition, sudden changes in color and brightness over several months might be indications of cometary activity, which has been speculated for a few objects. Filter # Filter Name Wavelength Description (nm) 1 Open (no filter) 2 U 365±66 Near Ultraviolet Light 3 B 445±94 Blue Light 4 V 551±88 Visual Light 5 R 658±138 Red Light 6 I 806±149 Near Infrared 7 Pinhole 8 Neutral Density Table 1: Filters of the CCD camera and descriptions

2 Making Measurements To make measurements you will need to acquire each night Calibration Files Standard star images Images of KBOs Calibration Files Calibration files will be necessary each night data is acquired. The images are used to remove obscuring dust on the optics, light sensitivity variations and additional random noise from the CCD camera. Improperly calibrating the images can cause results with larger error. Explained below are three kinds of calibration images which will be used. The bias frame is an exposure of zero exposure time. The bias frame will measure the amount of noise in each image. The electronics of the CCD camera will cause each pixel to register some amount of charge. The charge can then be misinterpreted as signal if not accounted for with the bias frame. The dark frame is an image of the same exposure time as the object, but receives no light. At a given exposure time a CCD chip will accumulate a certain amount of signal that is not related to an image. That is, if you repeat the same exposure under the same circumstances (temperature & time) the signal can be repeated. This means that if we take a dark frame and then take an image at the same temperature and for the same time as the dark frame the unwanted signal that is accumulated and recorded during the exposure on the dark frame will be on the image as well. So, if we subtract the dark frame from the image we can remove all the unwanted signal. Flat fields are images of evenly illuminated surfaces such as the twilight sky. Ideally, we would like to use flat field images taken that evening, however in case of poor weather conditions at sunset, you might use flat field images from a previous session. The images are useful in measuring the amount of dust or dirt in the optical system. A speck of dust on the mirror will be out of focus and appear on the image as faint a doughnut shaped region. Flat field images also take into account vignetting and other light varying problems. To remove these imperfections we divide the final astronomical image by the flat field. Standard Stars Standard stars are stars which have well known magnitudes in several filters. Standard star measurements will allow you to calibrate the observations of the KBO to the magnitude system. Use a standard star which is not too far from the KBO (< 5 ), at about the same airmass, and probably brighter than the KBO. Alternate observations of the standard star with the KBO, making sure you have at least three sets of standard star measurements and several images per filter in each set. Obtain measurements of the standard star in each filter used for the KBO. Some bright standard stars are listed in section H of the The Astronomical Almanac, provided in the control room. Use standard stars that are calibrated for the Johnson UBVRI filter system. You may also wish to use Landolt standard stars, which typically contain several stars in a given field. Since the Landolt standards are fainter than those given in The Astronomical Almanac, you will need longer integration times to get a good number of counts. The Landolt standards are located next to The Astronomical Almanac as a clipped stack of papers. It can also be found in the July 1992 issue of the Astronomical Journal (Landolt 1992).

3 Kuiper Belt Objects The KBOs have been broken into three main categories, Plutinos, classical objects and scattered disk objects (SDO). Plutinos share a 3:2 resonance with Neptune, similar to that of Pluto and make up roughly 25% of the known KBOs. Classical objects are those which do not have the 3:2 resonances and are found far enough from Neptune to avoid gravitational disturbances. Most of the classical objects are found between 42 and 48 A.U. and they tend to have small eccentricities and can have inclinations in excess of 30 (1996 RQ 20 and 1997 RX 9 ) to the plane of the solar system. By far, classical objects make up the largest collection of KBOs. Scattered disk objects have large eccentricities and inclinations. Perihelion distances of these objects is around 35 A.U., and due to perturbations by Neptune, they have aphelion distances that can reach 200 A.U. (1999 CF 119 ). Scattered disk objects only account for 3-4% of the known objects, but because of observational biases this is only a lower limit. Recently a new class of KBOs which includes the 3 categories above, as been discovered. Up to 8 binary KBOs (not including Pluto and Charon) have been detected to date. Since observational biases also limit the detection of binary KBOs, a lower limit of ~1% can be placed on the population of binaries. It is useful to check for the latest news about recently discovered objects, since new KBOs are reported on a regular basis. Even the brightest objects will be difficult to observe. For the observations you will require many moonless nights (or at least the moon can not be up during observing), 5-10 minute integration times, and the sharpest seeing. Even with this it will be difficult to identify the KBO from a faint star. If the seeing gets much beyond 4 arcseconds, the data will not be too useful. Constantly check the seeing, and adjust the focus every minutes. Since your eyes can easily pick out objects which move with respect to the background stars, you will need to animate the images. To determine the location of a KBO at a given time, you will need to generate an ephemeris. This can be done by entering the required information at the JPL ephemeris generator located at Once you move the telescope to that field, you can estimate where the KBO is in the field by comparing a field image from the telescope with the same field from the Aladin sky atlas. You can get to the Aladin sky atlas from Simbad at and clicking on the Aladin link at the top of the page. Table 2 lists several KBOs bright enough to observe from Braeside. Some of these objects have been studied using other instrument. References to previous results are given in the table. You can find these papers, and others, at Name V Opposition (2004) Class Reference 2002 AW Jan TNO M AZ Jan TNO 1999 DE Feb SDO JL CR Feb SDO 2003 FY Mar TNO 2002 JR Apr SDO Huya (2000 EB197) 19.6 Apr TNO JL GN Apr TNO SJ FX Apr SDO 2002 KX May TNO Ixion (2001 KX76) 19.7 May TNO M03 Quaoar (2002 LM60) 19.2 May TNO M03, O MS Jun TNO 2002 MS Jun TNO 2002 TX Sep TNO 1999 TC Sep TNO P02, D TD Oct SDO C02, C03, R UX Oct TNO 1995 SM Oct TNO D VS Oct TNO 2001 UO Oct TNO 2002 VE Nov TNO 2003 WT Nov SDO Varuna (2000 WR106) 20.1 Dec TNO JnS02, J01, L02 Table 2: Table of bright KBOs.

4 Data Acquisition To control Braeside from ASU, you will use the program XBOBS (X-based Braeside OBserving System). The program was specifically written to operate the telescope, dome and camera. You can view the telescope via the internet at To observe from ASU and control the observatory, do the following Log into the Braeside computer located in PSH-563-A2. Login name: observer, password: braeside Open an xterm by typing xterm & at a Konsole prompt. At the observer prompt, type in idl. At the idl prompt, type in xbobs In the region marked Telescope Control go to the Dome/Tel section, hit the pulldown button marked Select and hit the Auto Open selection. This will initialize the telescope and set it up for observing. You may watch and listen to the telescope as it sets up by turning the room lights and audio on in the region marked Dome Control. When the telescope is ready for you to use, the top row of red and green lights should be all green, and any warning boxes should have disappeared. Remember to turn the room lights off before taking data. In the region marked Target Acquisition and Pointing turn the Pointing Updates on. The boxes below the switch will fill in with numbers indicating the location of the telescope and where it is pointed. In the Dome Control panel, link the dome by selecting the Link option near the Dome buttons. If this is not done then the dome will not track with the telescope. Find the location of a bright star with in one hour east or west of the meridian at the time of observing. There is a catalog of bright stars built into the program. In the same section of the widget, find the region marked Catalogs and press the button marked Local. In a few seconds a window will appear. Scroll down to find the name of the star near the meridian. Highlight the star. Hit select. The information should appear the in boxes for the Destination. Then hit exit. Move the telescope to the destination by pressing the button marked Go There. You will be able to hear the telescope move into place. Once at the location, enter the observers names in the Observer box and press the Set button all located in the Observing Interface section. Then take a short exposure (~0.5s) of the region. Enter the exposure time in the box marked Exptime, and take the exposure by pressing the button marked Expose. The image will appear in the window. You want the star as close to the center of the frame as possible. If the star is not in the frame, take a longer exposure (~15s). If the star is bright enough (>4th magnitude) then spikes of star light should enter the frame. Move the camera in small increments by going to the Manual Positioning and pressing the N, S, E, or W buttons, and adjusting the increment in the center. You can move the telescope between arcseconds. Once the star is in the frame, you can center it by pressing the Recenter Image button in the Target Acquisition and Pointing section, then click on the center of the star, and then hit Go There. You will want to reset the coordinates by pressing the Reset Curr to Dest button. If the star looks like a doughnut, then you will need to focus the telescope. Go to Focus in the Observing Interface panel and adjust the focus by short amounts, press Set and take short exposures in between changing the focus. Once the star looks in focus, check the seeing estimate of the field stars by pressing Seeing Estimate in the image display panel and clicking on the center of a star. Typical seeing with Braeside is about arcsec, excellent seeing is 2.5 arcseconds and poor seeing is >~6 arcsec. While you adjust the focus, the seeing will decrease until you are at the best focus possible. Continue to get estimates of the seeing while you are observing. You can record the focus and seeing estimate on the observers log sheet in the comments column (Table 4). Once this is all set, you may go to your target. Enter the coordinates of the target and object name in the RA, Dec, and Object Name boxes in the Destination section each time you go the star, or you may add them to the local catalog. To add a star, open the local catalog, press the Add button, enter the information and then press the Add button in the new window. Move the telescope to the target as described above. Center the telescope on the target, set the filter you will be using and obtain an exposure. Get an image where the counts inside the target are between 50,000-60,000 counts in the raw image. You can check the number of counts by hitting the Track button in the image window panel. The mouse position (x and y) and counts (flux) will be displayed while you are moving the mouse over the image. Click the right mouse button to automatically move the mouse to brightest local pixel. You must click the left mouse button to deactivate tracking before proceeding. If the counts for the target are too low, for example 10,000 in 5 seconds, you can use the fact CCDs respond linearly, that is, if you want 60,000 counts, then extend your exposure time by a factor of 6. As the counts exceed

5 60,000, the response of the CCD changes from linear to logarithmic. Avoid using data with counts greater than 60,000. Images where the counts reach 65,536 are considered saturated and are usually notable by vertical bleeding of the stars. Finding a good exposure time may take several trials. You want images with as many counts as possible to be able to apply Poisson statistics for error analysis. Once you have the image set up the way you would like it, go to the Camera Control panel and turn on the Auto Increment option. When this is set to Off then the filename displayed in the Filename to be written box inside the Observing Interface panel will never change and you will write over the same image. Once the Auto Increment is set to On, then each exposure will save an image. When you have finished using the telescope for the night, close the telescope by selecting the Auto Close option in the Dome/Tel button in the Telescope Control panel. Use lights and camera to confirm the closure and parking of the telescope. Data Reduction The data should be analyzed using a language which you are comfortable with, such as C, Fortran, or IDL. Each file made at the telescope is written as a.fits files, and so you will need a language capable of reading the file in. The first section of each file is a header containing the observation information (filter, exposure time, RA and Dec, airmass, etc). To begin the reduction, subtract the bias image (B) from the data image (I(raw)). Then subtract the product of the dark image (D, which is normalized to one second) and the exposure time (t) of the image. Then divide the flat image (F). I ( clean) = I ( raw) B ( D t) Measure the light from the KBO, standard star and a star near the KBO (field star) by summing up the pixels within a circular region around the star. Since the light from an object spreads out like a 2-dimensional gaussian, make sure you have used an area large enough to include fainter pixels. Be very careful not to make the aperture too large for the KBO, because the KBO is so faint, you want to avoid including too much background. Use the field star to show that variations you may observe in the KBO are due to the KBO and not related to the weather conditions. The total number of counts for a object is an instrumental flux which can be converted to a magnitude using the following equation F m = 2.5 log10 where f is the instrumental flux per second, and m is the instrumental magnitude. The error in the flux will be assumed to follow Poisson statistics, that is ( f ) σ f = f which can be used to obtain the error in the magnitude as σ m = σ f To convert the measurements from instrumental magnitudes to an actual magnitude and correct for observing at different airmasses, you will need to make a plot of the standard star magnitude vs. airmass. Airmass is the is the secant of the angle between the object and the zenith. This value is stored in the header of each file. There is a linear relationship between magnitude and airmass. Each magnitude reported in literature is given at an airmass of 1, that is if the object was at the zenith. The slope of the plot will determine the relationship between magnitude and airmass. The offset between the instrumental magnitude and the accepted magnitude is the y-intercept. Apply the slope and y-intercept to the KBO to obtain magnitudes at 1 airmass. Make a plot of magnitude vs. time for the KBO. Phase the data to the period (which you may need to determine) to better show the periodicity of the object.

6 Start Time: End Time: Weather Conditions: Image Number Extension Number Filter Object Exposure Time Comments R Quaoar 450s Focus = -330, Temp=55 F Table 4: Observers log sheet

7 References Choi et al.(2002) [C02] Choi, Y.J., Prialnik, N., Brosch, N., 2002, Asain-Pacific Regional Meeting, p Choi et al. (2003) [C03] Choi, Y. J., Brosch, N., Prialnik, D., 2003, Icarus, 165, 101. Delsanti et al. (2001) [D01] Delsanti,A.C., Boehnhardt, H., Barrera, L., Meech, K.L., Sekiguchi, T., Hainaut, O.R., 2001, Astron & Astrophy., 380, 347. Dotto et al. (2003) [D03] Dotto, E., Barucci, M. A., Boehnardt, H., Rommon, J., Doressoundiram, A., Peixinho, N., de Bergh, C., Lazzarin, M., 2003, Icarus, 162, 408. Jewiti et al. (2001) [J01] Jewitt, D., Aussel, H., Evan, A., 2001, Nature, 411, 466. Jewitt & Luu (2001) [JL01] Jewitt, D.C. and Luu, J.X., 2001, Astron. J. 122, Jewitt & Sheppard (2002) [JnS02] Jewitt, D.C. and Sheppard, S.S., 2002, Astron. J. 123, Lellouch et al. (2002) [L02] Lellouch, E., Moreno, R., Ortiz, J., Paubert, G., Doressoundiram, A., Peixinho, N., 2002, Astron & Astrophy, 391, Marchi et al. (2003) [M03] Marchi, S., Lazzarin, M., Magrin, S., Barbieri, C., 2003, Astron & Astrophy, 408L, 17. Margot et al. (2002) [M02] Margot, J.L., Trujillo, C., Brown, M.E., Bertoldi, F., 2002, AAS, DPS meeting #34, # Ortiz et al. (2003) [O03] Ortiz, J. L., Gutieruz, P. I., Sota, A., Casanova, V., Teixeira, V. R., 2003, Astron & Astrophy, 409L, 13. Peixinho et al. (2002) [P02] Peixinho, N., Doressoundiram, A., Romon-Martin, J., 2002, New Astron., 7, 359. Rousselot et al. (2003) [R03] Rousselot, P., Petit, J.-M., Poulet, F., Lacerda, P., Ortiz, J., 2003, Astro & Astrophy, 407, Sheppard & Jewitt (2002) [SJ02] Sheppard, S.S., Jewitt, D.C., 2002, Astron. J., 124, 1757.