PROJECT 10 SUPERNOVA REMNANTS Objective: The purpose of this exercise is also twofold. The first one is to gain further experience with the analysis of narrow band images (as in the case of planetary nebulae) and the second is to examine the nature of the emission from an extended object. Emission observed in H α at 6563 Å and the forbidden lines of singly ionized sulphur at 6716 and 6731 Å may originate from photo ionization or from shock heated gas. Their relation can be a useful diagnostic tool. Observations: Numerous observations of supernova remnants have been performed so far at Skinakas Observatory resulting in a significant number of refereed publications in widely accepted journals. Newly detected structures have been proposed as candidate supernova remnants; optical emission has been detected for the first time from a number of already known radio remnants, while deeper observations have been performed in certain cases. The excellent site conditions, the 30 cm wide field telescope and the spectral capabilities of the 130 cm telescope offer unique opportunities for the study of this class of objects. The students will make use of one of the telescopes available at Skinakas Observatory to acquire images in the hydrogen and nitrogen emission lines H α [NII] and sulphur [SII]. Images of standard stars for the calibration of the narrow band images are also required. Flat field frames, 4 in each filter (H α +[NII], [SII] and continuum filter) Bias frames, a total of 10 spanning over the night Object frames, 1 at each filter (minimum). e.g. for G65.3+5.7 Standard star frames, 10, at least, at 3 different airmasses
Theory topics: Supernova remnants, narrow band filters, forbidden line emission, Balmer line emission. Analysis: The first step is to properly calibrate the target object images. Initially, the object images must be corrected for the bias level of the CCD camera. Subsequently, they will be corrected for the uniformity of the CCD by combining the individual flat field frames to create the master flat field in each of the available filters. Once corrected, the background level must be determined and subtracted. Depending on the complexity of the field of the object, a fixed value for the background may be sufficient. Otherwise, two dimensional images of the background should be created and subtracted. In the final step, the continuum image will be subtracted from the narrow band object images. However, this requires scaling the continuum image through photometry of field stars in order to calculate the scaling factor. Fig. 1. A raw image of an unknown structure which is believed to be a supernova remnant.
Contents: Supernova Remnants 1. Introduction 2. Processing 3. A few details 1. INTRODUCTION Supernovae may be distinguished into two major categories, type I and type II. Type I do not display hydrogen emission in their spectra. This category is further divided into Ia, Ib and Ic types depending on the presence of the element of silicon in their spectra. Type Ia supernovae display silicon in their spectra, while type Ib do not. If helium is also present in the spectra, then the type is considered to be Ic. Type II supernovae do show hydrogen emission in their spectra and are believed to originate from heavy stars (initial mass greater than 8 solar masses) that have exhausted their nuclear fuel. Type I supernovae are believed to originate from low mass stars (white dwarf) that accrete mass from a companion. When a certain density limit is exceeded the white dwarf will detonate through uncontrolled runway fusion of carbon and oxygen. During a supernova explosion vast amounts of energy (10 51 ergs) are released into the interstellar medium (ISM). The ejecta initially travel with velocities of the order of 10 4 km/sec and the shock front that has developed heats the ISM to temperatures around 10 7 K or more. This is the free expansion phase which lasts typically around 10 2 years. During this time the shock front has swept up mass equal to that of the ejecta, depending on the density of the local ISM. The temperature inside the remnant is roughly constant as well as the expanding velocity. The next phase is known as the adiabatic or the Sedov Taylor phase where the shock gradually decelerates and the ejecta mix with the shock material. The Sedov Taylor solution shows that the shell radius, R, scales as t 2/5, while the shock temperature drops as R 3. This phase typically lasts for 10 4 years. After this phase a thin but dense shell begins to form which encloses the interior gas found at temperatures of the order of million K. This phase is known as the pressure driven snow plow phase. The electrons begin to recombine with the atoms, thus releasing energy which leads to shell cooling. As the shell cools, more atoms are able to recombine and more energy is lost. The outward expansion slows down and after around 10 6 years all SNR material will be mixed up into the ISM enriching it with heavy elements. Currently, the types of SNRs may be distinguished in shell, composite and mixed morphology. During the various phases of SNR evolution emission is observed in X rays, UV, optical, radio and IR wavelengths. Here we are interested in the optical emission from such objects and to a first approach we will examine their H α and sulphur line radiation. In general, the optical emission from supernova remnants is the result of collisional
ionization and not photoionization. The students should note that the electrons that collisionally excite the atoms are energized by collisional ionization, while in planetary nebulae the electrons are energized by photoionization. The Table below lists some of the major emission lines observed in supernova remnants. Line Wavelength (Å) [ΟΙΙ] 3727 [ΟΙΙΙ] 4363 H β 4861 [ΟΙΙΙ] 5007 [ΝΙΙ] 5755 [ΟΙ] 6300 [ΟΙ] 6364 [ΝΙΙ] 6548 H α 6563 [ΝΙΙ] 6584 [SII] 6716 [SII] 6731 As we have mentioned in the beginning we are interested in line images of an object including the sulphur lines at 6716 and 6731 Å and the hydrogen line of H α. The forbidden [SII] line emission occurs in a region further behind the shock front where the temperature has dropped to around 10 4 K, while Balmer line emission at 6563 Å occurs closer to the shock front. Generally, the degree of ionization in supernova remnants behind the shock front is lower compared to planetary nebulae. The forbidden sulphur emission is thus stronger in SNRs and a criterion has developed to distinguish between remnants and HII regions or planetary nebulae. It is likely that if the ratio of [SII] to Hα is greater than around 0.4 then we are dealing with emission from a supernova remnant (Mathewson and Clark ApJ 178, L105, 1972; Sabbadin and D Odorico, A&A 49, 119, 1976; Vanessa et al. ApJ 118, 2775, 1999). Additional data are required in cases of doubt (e.g. non thermal radio emission, filamentary morphology, radio optical spatial correlation). 2. PROCESSING Arrange your observation time table so that a number of standard stars, listed in Table I, are observed in a range of airmasses. Try to include the airmass of your target in the corresponding range of your standard stars. Select the exposure times in a way that the photon counting errors are small without overexposing the CCD camera. You must also take care that the target observations are performed sequentially in order to avoid aligning the images at this stage. It is clear that in order to obtain proper results the target images must be projected to a common origin on the sky before combining them. This process will be described in detail elsewhere.
Table I Standard Star Properties Standard Star HR9087 HR718 HR5501 HR7596 HR7950 V magnitude 5.12 4.28 5.68 5.62 3.78 Right ascension 00 01 49 02 28 10 14 45 30 19 54 45 20 47 41 Declination 3 01 39 08 27 36 00 43 03 00 16 25 9 29 45 Integral I s (Ηα+[ΝΙΙ}) 6.4E 6 1.3E 5 3.7E 6 4.3E 6 2.2E 5 Integral I s ([SΙΙ]) 2.5Ε 6 5.4Ε 6 1.5Ε 6 1.8Ε 6 9.1Ε 6 Fig. 2. A combined flat field frame in [SII] 6720 Å.
Fig. 3. Bias frame showing a uniform distribution. Narrow band or interference filter images are, initially, processed in a similar manner like broadband images. However, this is correct up to the stage where instrumental magnitudes are obtained. The processing after this step differs significantly. Let's see in more detail the analysis of the spectrophotometric standard stars. The interference filters are characterised by full width half maxima of ~10 20 Å and it is reasonable to expect that the atmospheric absorption, the quantum efficiency of the CCD and the response of the telescope optics are basically constant over the band pass. Then, the light recorded by the CCD, after passing through the Earth s atmosphere and the telescope optical system, can be expressed as where 0.4kχ t exp N( adu) = Ao 10 N star ( λ) TF ( λ) λdλ hc (1) N star, is the total source counts, A o, is the zero point of the magnitude scale, K, is the extinction due to our atmosphere, χ, is the airmass of the star during the observation, t exp, is the exposure time in sec, h, is Planck's constant in CGS units, c, is the speed of light in CGS units, T F (λ), is the transmission of the filter as a function of wavelength, N star (λ), is the known energy distribution of the standard star in erg /sec/ cm 2 /Å.
Reducing this relation to magnitudes and rearranging terms, we obtain N( adu) 1 2.5log( ) + 2.5log( N star ( λ) TF ( λ) λdλ) = 2.5log( Ao ) + kχ t hc exp (2) The measured parameters are the instrumental magnitude, the airmass during the observation and the exposure time in seconds. The value of the integral depends on the specific standard star and filter. These are supplied for a number of bright stars and filters in Table I and for conditions typical to Skinakas Observatory. It is evident that the corrected magnitude varies linearly with the airmass and a linear fit of the left hand side of eq. (2) vs. airmass would determine the slope k and the y intercept 2.5 log(a o ). You can write your own program to perform the fit or use any spreadsheet program with fitting capabilities. Once k and A o are determined, you can proceed to calibrate the object data (SNR). In the case of an unknown star, we can only calculate its integrated flux by reverting eq. (1) and calculating the integral part, since A o and k are now known. This is what would be measured if the telescope was placed at the top of the Earth's atmosphere. The object data are also processed in the usual way (bias, flat field, etc.) but in addition, the sky background must be subtracted. In order to perform this action, select several areas in your image which are free from nebular emission and measure the sky background. Write down and examine these intensities along with the standard deviation given by the software. Use an average value if the values agree within the errors, otherwise the construction of a two dimensional sky background image should be considered. Once the sky level has been subtracted, the calibration procedure can be applied. However, equation (1) cannot be used because we are now dealing with an extended object and its flux should be calculated per unit sky projected area. The slightly modified form of equation (1) reads N ij t ( adu) δ hc 0.4kχ exp 2 = Ao10 N ij ( λo ) TF ( λo ) (3) It was assumed that any emission line behaves like a delta, δ(x), function at λ o and the integral is simply evaluated at this wavelength measured in Angstrom. The indices ij refer to pixel coordinates in the image where the observed intensity is N ij in adu. A o and k are already determined through the calibration process. The airmass χ is a weighted average because during long exposures the actual airmass may vary. You can use the following formula for the effective airmass:
4 end χ start + χ middle + χ χ = (4) 6 The exposure time is typically several hundreds seconds since supernova remnants are faint objects, while δ is the pixel size of the CCD in arcseconds. Consult the Skinakas operator for the actual value of δ used during your observation. If the camera pixels are not squared, then δ 2 should be replaced by the product δ x δ y. The transmittance of the filter at the specified wavelength λ o of the emission line is T F (λ o ). Equation (3) can be solved for the unknown flux N ij (λ o ) which will be a two dimensional image in units of erg/sec/cm 2 /arcsec 2. It describes the flux from the target object at wavelength λ o that we would observe if the telescope was located at the top of the Earth's atmosphere. If there is more than one image available in the same filter, the calibration procedure must be applied to each one of them. Use of the H α +[ΝΙΙ] filter allows for the immediate application of Eq. (3) by solving for Ν ij (λ ο =6563 Å). It is proposed to use two narrow sulphur filters for the observations. The wide 6720 Å filter intercepts both sulphur line but at different transmittances. Typically, we may write: N ij (6563 Å) = N ij (adu) 0.56 hc/t exp 10 0.4 k χ /A o /δ 2, and N ij (6716/6731 Å)= N ij (adu) 2.9 hc/t exp 10 0.4 k χ /A o /δ 2, In this way, the two line images at 6563 Å and 6716+6731 Å have been estimated separately and can be used for measurements. They will identify the same areas of interest in all line images and measure with the aid of the software the corresponding flux. For any given area, the flux at the 6716+6731 Å image will be measured and subsequently, divided by the flux measured at the 6563 Å image. The ratio of the two values will be the ratio [SII]/H α. Summarizing, it is proposed to: 1. Determine the bias level (mean and standard deviation) 2. Subtract the mean bias level from all available frames 3. Create the master flat field frame for all different filter used in this study 4. Divide all available data frames with the corresponding flat field frame 5. Reduce the standard star data to determine the calibration factors 6. Apply the calibration to the target object frames 7. Measure the fluxes and determine the ratio of [SII]/H α
3. A FEW DETAILS Case 1: Several images are available for a single filter or a 2D image of [SII]/H α is necessary. If more than one image is available for any of the filters, then these must be projected to a common origin on the sky. Alternatively, one of the images can be used as a reference. This exercise does not address the problem of the alignment of images but a few hints are provided. Use the image with the minimum FWHM of field stars as the reference image. Smooth all images with a Gaussian filter to match the image with the worst FWHM. Images should be averaged after their calibration Case 2: In a next level of training, students may wish to subtract the continuum image from each of the line images. This means that the images must be properly aligned. Once aligned identify stars common to all images and measure their fluxes. Calculate the ratio between the line image fluxes and the continuum fluxes, check the mean of all values along with the standard deviation and omit any variable stars from the list. Perform the subtraction with the scale factor you have just calculated and smooth the resulting image to suppress any residuals. The final image can be used for further scientific analysis. Fig 4. The continuum subtracted image of filamentary structures are seen on the left image, while the original calibrated image is seen on the right.