Development of Fast Imaging Solar Spectrograph and Observation of the Solar Chromosphere

Transcription

1 A Dissertation for the Degree of Doctor of Philosophy Development of Fast Imaging Solar Spectrograph and Observation of the Solar Chromosphere Department of Astronomy and Space Science Graduate School Chungnam National University By Hyung-Min Park Advisor Titular Advisor Jongchul Chae Yu Yi February 2011

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3 Development of Fast Imaging Solar Spectrograph and Observation of the Solar Chromosphere Advisor Jongchul Chae Co-advisor Kyung-Suk Cho Titular Advisor Yu Yi Submitted to the Graduate School in Partial Fulfillment of the Requirements for the Degree of Doctor of Philosophy October, 2010 Department of Astronomy and Space Science Graduate School Chungnam National University By Hyung-Min Park

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5 To Approve the Submitted Dissertation for the Degree of Doctor of Philosophy By Hyung-Min Park Title: Development of Fast Imaging Solar Spectrograph and Observation of the Solar Chromosphere December 2010 Committee Chair Dr. Kap-Soo Oh Chungnam National University Committee Dr. Yu Yi Chungnam National University Committee Committee Committee Dr. Soo-Chang Rey Chungnam National University Dr. Jongchul Chae Seoul National University Dr. Kyung-Suk Cho Korea Astronomy and Space Science Institute Graduate School Chungnam National University

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7 Abstract Traditionally, solar observations have been performed by ground-based instruments. Next to the photosphere, the solar chromosphere has been studied well for a long time. It is well known from high-resolution observations that chromospheric features are fine structured, short lived, and dynamic. In obtaining physical parameters, spectrograph-based observations are more effective than filter-based observations. Through imaging spectroscopy using a spectrograph, chromospheric features and dynamics can be revealed. The biggest telescope, New Solar Telescope (NST), was recently built at Big Bear Solar Observatory. NST has a capability of high spatial resolution, 0.08 at 500nm, with the aid of Adaptive Optics. As one post-focus instrument of NST, an imaging spectrograph, called Fast Imaging Solar Spectrograph (FISS) was proposed and constructed by Korean researchers to study the solar chromosphere. This thesis mainly describes our contribution to the development of this spectrograph and early results. FISS is a grating-based spectrograph with high spectral resolution, high time cadence, and the capability of imaging. It has a mount of Littrow type, records dual bands simultaneously, and uses an Echelle grating as the disperser and performs imaging using a field scanner. We describe its optical design and performance estimation. Software development, construction and integration of each component were completed in Korea Astronomy and Space Science Institute. Through tests, we confirmed that the performance of the spectrograph has come close to our expectation. After FISS was installed on the vertical table on the Coudé room at Big Bear Solar Observatory, we observed various chromospheric features: active regions, quiet regions, filaments, prominences and so on. We determined physical parameters of limb prominences observed by FISS. By applying a non-linear least square fitting of a radiative transfer model to the profiles of Hα line and CaII 8542Å line, we derived physical parameters of the prominences. The ranges of temperature and non-thermal velocity are found to be 7,500-13,000 K and 5-11km/s, respectively. The maximum 1

8 temperature of prominences is found to be below 20,000 K. It is expected that FISS will contribute to revealing fine structures and the dynamics of the solar chromosphere with high resolution. 2

9 Contents 1 Introduction Background New Solar Telescope Adaptive Optics Necessity of a New Instrument Outline Instrument Introduction Design Concepts Optical design Estimation of spectral parameters Theory Estimation of parameters Intensity Estimation Components Scanner Slit Collimating/Imaging Mirror Disperser Filter CCD Cameras

10 2.7 Integration Hardware Software Installation Test Configuration of feed optics Laser Test Sunlight Test Conclusion Early Observations Data processing Sunspots Filaments Active Regions Quiet Regions Prominence Conclusion Determination of physical parameters of prominences Introduction Model of Radiative Transfer Data and Analysis Results Conclusion Conclusion 103 Bibliography 105 4

11 List of Figures 1.1 Perspective drawing of New Solar Telescope (courtesy by BBSO website) New Solar Telescope. A red square represents the Nasmyth bench Speckle reconstruction image of a sunspot (AR1084) with AO on July 2, 2010 at TiO (706nm) band The concept of AO Optical ray design of FISS by using ZEMAX Spot diagrams of Hα and CaII8642 by using ZEMAX A perspective drawing of FISS The field scanner The slit A collimator/imaging mirror An Echelle grating Filter transmission of Hα (left) and CaII (right). Dashed lines represent the wavelength at Hα and CaII8542, respectively A filter adapter DV887/DV885 EMCCD Quantum efficiency curves of DV887 (upper) and DV885 (lower) A layout between hardware connections. In this picture, M means a motor to move a CCD mount A control program of FISS

12 2.14 The position of FISS components Installation of FISS on vertical table in BBSO AO (left) and FISS (right) in Coudé room Feed optics with the laser. It contains laser, telescope as the beam collimator, iris, and convex lens The primary mirror of the coelostat Enlarged laser spectrogram (upper) and its horizontal profile (lower). In laser spectrogram, x-direction and y-direction indicate spectral domain and spectral domain, respectively. In lower plot, a dashed line presents gaussian fitting Efficiency curves of echelle grating at Hα (left) and CaII8542 (right) The solar spectrogram (left) and its line profile (right) for Hα Two spectrograms of Hα (left) and CaII8542 (right) at the same time Raster scan image of the resolution panel on May 10, On right side (square), the distortion appears in vertical direction because of the fluctuation caused by an airplane Raster scan images of the Sun constructed at Hα (left) and CaII8542 (right) on December 14, Horizontal lines on CaII8542 image are caused by the CCD anomaly at IR Examples of raw spectrogram (left top), slit pattern (right top), processed spectrogram (left bottom), flat pattern (right bottom) at Hα The comparison of raw data profiles with compressed data profiles. In upper row, solid lines and red dashed lines represent raw data profiles and compressed data profiles at Hα and CaII8542, respectively. Differences between raw data profiles and compressed data profiles are seen in lower row

13 4.3 Scan images of a sunspot observed on June 30, FOV is Dashed lines represent a position of spectrograms in Figure Sunspot spectrograms of Hα and CaII Sunspot images on July 22, FOV is Speckle reconstruction image of sunspot group (AR 11089) at TiO band (706nm) on July 22, Scan images of quiescent filament on July 29, FOV is Scan images of a quiescent filament on July 22, FOV is Scan images of a quiescent filament on July 16, FOV is Scan images of an active region filament on June 25, FOV is Emerging flux region associate with active region on July 22, FOV is Scan images of quiet region near disk center of the Sun on June 25, FOV is Hα (upper row) and CaII (lower row) spectrograms taken at the solar limb and a prominence outside it Hα full disk image on June 30, 2010 in BBSO. Two squares represent field of view of FISS Scan images at Hα, CaII in east (left column) and west limb prominence (right column) A raw spectrogram (upper) and a correction spectrum subtracted aureole light (lower) at Hα Two points at the east prominence of Hα (upper) and CaII8542 (lower)

14 5.5 Fitting results of Hα and CaII8542 on position A (upper row) and B (lower row). Left column represents Hα profiles and fittings, while right column shows CaII8542 profiles and fittings Fitting results at two positions with τ 0 -variation. Black, green, red dashed line represent a spectrum profile from data and fitting results with τ 0 of 2.5, 0.3, respectively. Blue dashed lines present τ 0 as free parameters Fitting results in Hα and CaII8542 at both the position A (upper) and the position B (lower) Fitting results in Hα and CaII8542 at both the position A (upper) and the position B (lower) Positions of interest with height in the east prominence Temperature variations with height Sampling positions of Hα (upper) and CaII (lower) in the east prominence Doppler width ( λ D(Hα), λ D(Ca) ), temperature, and nonthermal velocity plot (left) and temperature non-thermal velocity plot (right) in the east prominence Sampling positions in the west prominence Doppler width ( λ D(Hα), λ D(Ca) ), temperature, and nonthermal velocity plot (left) and temperature non-thermal velocity plot (right) in the west prominence The position where the line broadening is biggest at Hα (upper) and CaII8542 (lower)

15 List of Tables 1.1 NST specification Base requirements of FISS Input parameters for optical design of FISS Input parameters to estimate the performance of FISS Estimation values of FISS throughput Parameters and estimation values to calculate photon counts Specifications of CCDs Typical observing parameters

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17 Chapter 1 Introduction 1.1 Background The Sun is the closest star from us and it s only a star whose surface we can see in detail. The violent activity of the Sun, like flares or CMEs and so on, may affect the Earth seriously, so the monitoring of solar activity becomes important. Phenomena in the solar photosphere such as the variation of sunspot number have been studied well for a long time. On the other hand, features in the chromosphere could not be observed except during the total solar eclipse for a while. The routine observations of the chromosphere became possible only later with the help of high spectral resolution observations using either narrowband filters or spectrographs. What is really annoying in ground observation is the troublesome atmospheric seeing. It makes the shape of an object spread, so it is hard to achieve high angular resolution. Several ways have been suggested to overcome the seeing. First of all, selecting a good site is important. In case of nighttime observations, the site is located at a mountain of high altitude that is far away from cities to avoid artificial lights. In daytime at such a site, however, it is easy for the sunlight to heat the ground, so the atmospheric convection near 11

18 the ground is a serious problem. If a solar observatory is located on a lake, water keeps the ground cool even during daytime, and hence the air near the ground can be stable. In reality some solar observatories are located on lakes including Big Bear Solar Observatory, Udaipur Solar Observatory, Huairou Solar Observatory, and so on. Another way of overcoming the seeing is to use the instrument like adaptive optics to compensate for image degradation. With the use of adaptive optics, one can achieve diffraction-limited resolution. There are two kinds of solar observing system: filter-based ones and spectrograph-based ones. In case of filter-based systems, either Fabry-Perot or Lyot filters are mainly used to successively record light in different narrow spectral ranges. The main advantage of these systems is that they can take images of a large field of view quickly, but it takes much time to tune many wavelengths necessary for the construction of line profiles. Interferometric BIdimensional Spectrometer (IBIS) is a representative instrument (Cauzzi et al., 2008). By contrast, a spectrograph-based system can take much spectral information at a specific position, whereas it takes much time to scan a large field of view. Multi-Channel Subtractive Double Pass (MSDP) spectrograph, the CCD imaging spectrograph at the Solar Tower Telescope of Nanjing University, the Mees CCD (MCCD) imaging spectrograph, NRL High Resolution Telescope and Spectrograph (HRTS) onboard sounding rocket are representative (Mein, 1977; Mein et al, 1997; Ding et al., 1995; Penn et al., 1991; Dere et al., 1986). For a long time, solar observations of highest quality have been performed at Big Bear Solar Observatory (hereafter, BBSO). This observatory is located on the Big Bear Lake at high altitude (about 2000m), and now has the biggest solar telescope in the world. Adaptive optics is being applied to correct atmospheric fluctuations, and two instruments using Fabry-Perot etalons will be installed soon. These instruments are mainly intended to observe the solar photosphere. We are interested in the study of the chromosphere using our new spectrograph system that was recently installed at BBSO. 12

19 1.2 New Solar Telescope The observation system of BBSO consists of two components: New Solar Telescope (hereafter, NST) and post-focus system. Four post-focus instruments InfraRed Imaging Magnetograph (IRIM), Visible Imaging Magnetograph (VIM), Correlation Tracker (CT), Cryogenic IR Spectrograph and Filtergraph are already installed or will be installed. NST has two foci: Nasmyth and Coudé. The Nasmyth bench is located on the side of the telescope (see Figure 1.2). CT and Filtergraph have already been installed on the Nasmyth bench. Other instruments plan to be installed in the Coudé room inside one floor beneath the telescope. Now two instruments, Adaptive Optics System and FISS have been installed on each vertical table in the Coudé room. Light from the Sun is collected by the telescope and fed into the Adaptive Optics before FISS. NST has been installed to replace the old 26-inch telescope. NST is a 1.6 meter diameter, off-axis Gregorian system consisting of a parabolic primary mirror, and is currently the biggest solar telescope in the world. Details of NST can be found in Figure 1.1. In order to collect more light and reduce stray light, it was designed to be off-axis. Detailed specifications are given in Table 1.1. Although BBSO is located at a good site, it s hard to achieve diffraction limited images directly from NST because atmospheric disturbance still affects. Such images can be obtained with the aid of Adaptive Optics (hereafter, AO) system and speckle reconstruction process. Speckle reconstruction is used to high-resolution images from a burst of short-exposure raw images. Figure 1.3 presents a sunspot image obtained by applying AO and speckle reconstruction. This image is the most detailed ever obtained in visible light. AO will be described in the next Section briefly. 13

20 Figure 1.1: Perspective drawing of New Solar Telescope (courtesy by BBSO website) 14

21 Figure 1.2: New Solar Telescope. A red square represents the Nasmyth bench. 15

22 Figure 1.3: Speckle reconstruction image of a sunspot (AR1084) with AO on July 2, 2010 at TiO (706nm) band. 16

23 Table 1.1: NST specification parameters values clear aperture 1.6 m effective focal length 83.2 m wavelength range µm mounting equatorial mount with friction-based slew field of view diffraction limit 0.08 at 500 nm 1.3 Adaptive Optics Due to inhomogeneities of the atmosphere, image degradation always occurs. Adaptive Optics is the device of real-time image compensation. Typically an AO system has three components: wavefront sensor, deformable (adaptive) mirror, and control system (see Figure 1.4). The disturbance by the seeing is manifest in the offset positions of sub-images focused by individual lenses of the lenslet array. Based on sub-images, the control computer calculates the offsets and send a command to the deformable mirror to correct for them. Since the atmospheric seeing fluctuates fast, the feedback of AO should operate at high frequency. Meanwhile, in solar observations, contrary to nighttime observation, the object we want to see is extended and there is no point sources available as guides. Alternatively, objects which have high contrast against background are used as guide sources. For example, the granulation in quiet regions, sunspots and pores in active regions are good candidates for guide sources. AO in BBSO is used at two wavelengths, G-band (430.8nm) and TiO (706nm). Images are better corrected at TiO than G band as AO works better in IR wavelengths than visible wavelengths. 17

24 Figure 1.4: The concept of AO. 1.4 Necessity of a New Instrument As we mentioned above, initially planned post-focus instruments focus on the study of the solar photosphere. Thus a need of a new instrument for studying of chromospheric phenomena was recognized. Generally many people think that the chromosphere is just a thin layer which links the photosphere with the corona, and which is homogeneous and in a static state. However, it turns out that the chromosphere is inhomogeneous and dynamic. Chromospheric observations with high resolution reveal many dynamic features and fine-structures including spicules, mottles, jets, fibrils, filament/prominences and so on. These features and their dynamics are still poorly understood. Furthermore, such chromospheric phenomena affect the dynamics of the corona, so the study of the solar chromosphere is important. In most cases, the Sun is observed using narrow filter like Febry-Perot at one wavelength. In case of the photosphere, filter-based observations are suitable because photospheric lines are narrow and the shift of wavelength is small. 18

25 However, if we try to observe the solar chromosphere, a filter-based system is not suitable anymore because chromospheric lines are broader and line shift is larger than photospheric lines. For example, we observe violent events such as filament/prominence eruption, we can see the event while the line centers are within filter bandwidth. If the speed becomes faster, we may not see the event anymore since the line centers may be shifted beyond the bandwidth. Of course, the problem caused by narrow bandwidth can be overcome by tuning the filter wavelength in series. But this method limits the spectral information seriously. Thus we came to consider a grating-based spectrograph to see the solar chromosphere with high spectral resolution, high time resolution, and high spatial resolution by the aid of NST. 1.5 Outline This thesis describes our contribution to the development of the new instrument, Fast Imaging Solar Spectrograph, and some early scientific results from it. Chapter 2 deals with all things of our instrument from concept to installation and simple theory. Results of lab tests are listed in Chapter 3, and we will show early observation of the solar chromosphere in Chapter 4. The study of quiescent prominence observed using FISS is described in Chapter 5. 19

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27 Chapter 2 Instrument 2.1 Introduction The chromosphere, a red color sphere, is a thin layer just above the photosphere of the Sun. During a total solar eclipse, it can be seen as a thin and red color layer by a naked eye. It links with the photosphere and corona of the Sun. It is from the observation of solar chromosphere that hydrogen and helium were found to be the main chemical constituents. There are many conspicuous features in the solar chromosphere including filaments or prominences, plages, active regions, and spicules. The dynamics and structures of these features, however, are not fully understood even though the solar chromosphere has been observed for a long time. High resolution observations have indicated that many chromospheric features have fine structures that are usually dynamic. For example, filaments/prominences are known to consist of many thin threads ( km width). Threads are short-lived, and usually contain mass flows along them. To study the fine-scale structures and dynamics of the chromospheric features, we needs an observing instrument which should have high angular resolution, high spectral resolution and high temporal resolution. For high angular resolution, the telescope must have an aperture which is big enough to reveal fine chromospheric structures 21

28 in detail. NST, the biggest telescope in the world, has been constructed at BBSO. It has the highest spatial resolution we have ever seen, 0.08 at 5000Å. As we mentioned in Chapter 1, initially planned observing instruments focus on the study of the solar photosphere. Accordingly a new instrument to see chromosphere was needed. With this necessity, we have developed the instrument: the Fast Imaging Solar Spectrograph (FISS). 2.2 Design Concepts Our scientific goal is to understand physical characteristics of the solar chromosphere, focusing on fine-scale structures and their associated dynamics. For this we need a spectrograph which has high spatial resolution, high temporal resolution, and high spectral resolution. Using this instrument, we specifically want to determine physical parameters such as temperature, density, chemical composition, and so on. If we observe the Sun at two or more wavelengths in high spectral order simultaneously, we can estimate several parameters precisely. Moreover, we can do imaging by moving the solar image across the slit with fast speed. That is the basic idea of the instrument. To realize the idea, we have refined the notion of the spectrograph. Spectrograph concepts are as follows: The spectrograph has dual channels to see the solar chromosphere at two different wavelengths simultaneously. The standard lines to observe the solar chromosphere are the Hα line and CaII IR line at 8542Å. Other lines such as CaII H & K or HeI in NIR can be used in the single channel mode. A diffraction grating is used as the disperser of the spectrograph. By using this, we can get spectral data at high spectral resolution with a large free spectral range. 22

29 Table 2.1: Base requirements of FISS Parameters Name Spectrograph Imaging S/N signal to noise ratio > R resolving power > 10 5 > x spatial resolution across the slit < 0.2 < 0.2 y spatial resolution along the slit < 0.2 < 0.2 t temporal resolution 60s 10s N s number of step per scan A s scanned area(fov) The spectrograph has simple configuration and compact size. If the spectrograph has simple configuration and small size, it is easy to repair and maintain. The spectrograph has the capability to do imaging. By successively scanning the position of the slit over the Sun within a short time, it is possible to do imaging with high cadence. Although the quality of imaging is poorer than a filtergraph, physical information contained in the spectra is abundant enough. To take fast scan, the stabilization has to be ensured. Since NST uses the AO system to correct for the seeing, we don t have to care about atmospheric fluctuation. Meanwhile, a fast camera is essential to obtain a series of spectrograms at high cadence. Table 2.1 shows the requirements for observation. Based on them, we have refined the concept of spectrograph in detail. To realize general ideas, we drew up guidelines as follows: The mounting of FISS is of Littrow type. It allows the spectrograph to have a simple configuration. Since incident angle and diffraction angle are close to blaze angle, a single mirror/lens can play as both a collimator and an imager. 23

30 Table 2.2: Input parameters for optical design of FISS Name Parameter Value F-ratio of incident beam F i 26 focal length of incident beam f i 41.6m collimator/image focal length f col 1.5m grating groove density d 79/mm grating blaze angle ϕ 63.4 deflection angle for CCD A α β 0.93 deflection angle for CCD B α β 1.92 An Echelle grating is used as the disperser to get spectra of high spectral orders. It is suitable for taking the data which have high spectral resolution and wide spectral range. We intend to obtain the data with high spectral resolution at a specific wavelength band, so a bandpass filter is adopted to remove the order-overlapping. A field scanner is used. The scanner enables the solar image to move across the slit. With the interlocking between the field scanner and the detector, we can do imaging. It will operate two ways: imaging mode and spectrograph mode. FISS has been designed based on these guidelines. Design parameters are listed in Table Optical design The design has been performed using the optical design program, ZEMAX. The pair of wavelength, Hα and CaII8542 is considered. A paraboloid mirror is used as the collimator/imager. Figure 2.1 is the optical layout based on the design parameters. The light dispersed by the grating is reflected by the 24

31 Figure 2.1: Optical ray design of FISS by using ZEMAX paraboloid mirror, and goes into each of two detectors. Figure 2.2 represents the spot diagrams at Hα and CaII8542, respectively. The circle around each spot represents the diffraction limit at the wavelength and is larger than the spot formed by the incident beam. Since the diffraction limit size is proportional to the wavelength, the size of diffraction limit at CaII8542 is larger than the airy disk at Hα. 2.4 Estimation of spectral parameters Theory The diffraction theory starts with the grating equation mλ = d(sin α + sin β). (2.1) 25

32 Figure 2.2: Spot diagrams of Hα and CaII8642 by using ZEMAX 26

33 where m is spectral order, λ, wavelength, d, groove distance between adjacent grating grids, α, incidence angle, β, diffraction angle, respectively. In order to understand the spectrograph, spectral order should be calculated first. FISS uses Littrow mounting, that means incident angle, diffraction angle and blaze angle of grating are almost the same(α β ϕ). Using this relation, spectral order at specific wavelength can be calculated. The spectral order, m, is estimated by 2d sin ϕ m = R n ( ), (2.2) λ where ϕ is blaze angle, R n is the rounding off function that yields the integer close to the argument. We introduce the difference between incident angle and diffraction angle, θ = α β. Removing β in Equation 2.2 using this expression of θ, we obtain the non-linear equation for α sin α + sin(α ϕ) mλ d = 0. (2.3) This equation can be solved iteratively for α using the Newton-Rapson method. Linear Dispersion defines the extent to which a spectral interval is spread out. It represents the ability to resolve fine spectral in detail. The expression of linear dispersion is as follows: where f cam is the focal length of camera. Linear dispersion per pixel is given by dλ dx = d cos β mf cam, (2.4) δλ = d cos β mf cam w det, (2.5) where w det is a pixel size. Based on linear dispersion per pixel, we can estimate the spectral coverage on CCD chip. Spectral coverage of detector is given by dλ = d cos β mf cam W det, (2.6) 27

34 where W det is the size of detector. Grating resolution is given by where W grating is grating width. δλ grating = λd mw grating, (2.7) Spectral purity measures the degradation by finite slit width. Incident beam that goes through the slit of finite width contains a finite spread of the incidence angle and brings about the spread of diffraction angle so that resolution gets worse. If we want to reduce spectral purity, either the distance from slit to collimator should be long enough or slit width should be narrow enough. Spectral purity (spectral resolution of the slit) is as follows: δλ sp = cos α w s f col d m, (2.8) where w s is slit width, f col is the distance between slit and collimator. Detector resolution is given by δλ det = 2 δx dλ dx = 2 cos β δx f cam d m, (2.9) where δx is a pixel size of detector, f cam is the distance between imager and detector. The net spectral resolution before sampling is given by δλ D = (δλ grating ) 2 + (δλ sp ) 2, (2.10) and spectral resolution after sampling by δλ D = (δλ grating ) 2 + (δλ sp ) 2 + (δλ det ) 2. (2.11) The resolving power is dimensionless, representative value for the spectrograph. It is proportional to the smallest difference between specific wavelengths. The equation is given by R = λ δλ D, (2.12) where δλ D the smallest difference in wavelengths that can be just distinguished. 28

35 2.4.2 Estimation of parameters Using relations associated with diffraction theory, spectrograph parameters are estimated. Input parameters to calculate throughput are listed in Table 2.3. Table 2.4 shows spectrograph throughput. 2.5 Intensity Estimation Ahn et al. (2008) estimated the shapes of spectral profiles of Hα and CaII8542 lines and calculated photon counts on the CCD chips. The photon count at the continuum (N e ) is expressed as follows: N e = q τ atm τ ins I C λ hν λ π = q τ atm τ ins I C x y hc Ω λ A t (2.13) λ t, 4F 2 where ω is solid angle of the telescope, A, unit area, F, focal ratio of incident beam, respectively. Other parameters and values are listed in Table 2.5. The photon count in Hα and CaII8542 are , , respectively. If this value is true, Analog-to-digital converting unit(adu) count on CCD should be too high to detect all. However photon count obtained by tests was found to be much lower than these values. Furthermore, the values obtained at Korea Astronomy and Space Science Institute (hereafter, KASI) were found to be much different from those taken at BBSO. This discrepancy between calculation and determination is due to their wrong assumptionon the instrument transmission. Ahn et al. regarded the instrument transmission as We found that this value was much overestimated. The optical configuration for NST and AO system has lots of optical surfaces like mirrors and lens. They didn t consider the transmission of all the components, mirror, filter, grating and so on. Instrument transmission calculated again is found to be about Based on this new value, we obtain photon counts of and for Hα and CaII 8542, respectively, corresponding to to 29

36 Table 2.3: Input parameters to estimate the performance of FISS Name Parameter Value telescope diameter D 1.6 m F ratio F 26 focal length of the incident ray f 41.6 m slit width w 16 µm slit height h 10 mm collimator/image focal length f col 1.5 m grating groove density d 79/mm grating blaze angle ϕ 63.4 deflection angle for CCD A α β 0.93 deflection angle for CCD B α β 1.92 camera A pixel size x y 16 µm 16 µm camera B pixel size x y 8 µm 8 µm CCD A format N x N y CCD B format N x N y CCD A spectral order m 34th at Hα CCD B spectral order m 26th at CaII8542 scan step size x s 16 µm number of steps per scan N s 600 scan time t 60 s scanned area A

37 Table 2.4: Estimation values of FISS throughput Name Parameter Hα 6563Å CaII 8542Å incidence angle α diffraction angle β spectral order m linear dispersion dλ/dx 1.2 må/µm 1.6 må/µm spectral coverage per a pixel δλ c 19 må 26 må spectral coverage on chip dλ c 9.8 Å 12.9 Å grating resolution δλ grating 19.7 må må detector resolution δλ det 38.1 må 51.3 må spectral purity δλ sp 18.5 må 24.2 må spectral resolution before δλ 27 må 41 må spectral resolving power before λ/δλ 243, ,000 spectral resolution after δλ 46 må 48 må spectral resolving power after λ/δλ 141, ,000 beam width on grating W grating 124 mm beam height on grating 57.7 mm 31

38 Table 2.5: Parameters and estimation values to calculate photon counts Parameter Name Ha CaII8542 q quantum efficiency τ atm atmospheric transmissivity λ spectral coverage of a pixel (må) I C continuum intensity (erg/cm 2 /str/sec/å) N e photon count of continuum N data number (DN) τ ins instrumental transmissivity x physical width of a CCD Pixel (µm) 16 y Physical height of a CCD Pixel (µm) 16 t CCD exposure time (sec) 0.03 F focal ratio of incident beam ADU and 1052 ADU, respectively. These values are close to the values obtained by observations. To enhance instrument transmission, it is necessary to replace the aluminum mirror coating by the enhanced-silver because the reflexibility of silver is larger than aluminum. 2.6 Components Typically, a spectrograph is composed of a slit, a collimator, a disperser, an imager and a detector. In addition to these, FISS has a field scanner and bandpass filters. As we mentioned, a portion of paraboloid mirror serves as both collimator and imager. Figure 2.3 shows the FISS configuration. CCD A and CCD B are to take spectrograms at Hα and CaII8542, respectively. Each component is described in this section. 32

39 33 Figure 2.3: A perspective drawing of FISS

40 Figure 2.4: The field scanner Scanner Generally Littrow spectrograph uses rotating mirror in front of slit to sweep an image of luminous object like the Sun. But FISS does not use such a rotating mirror, and uses instead a field scanner which is composed of two flat mirrors and a linear stage motor as the movement device. This scanner does not require a pupil as the location, and hence it makes the optical design of the instrument much simple. The field scanner installed on the vertical table is shown in Figure 2.4. It uses JSMD linear motor by Justek and PMAC Mini board by Delta Tau. The field scanner displays two kinds of motion: step-by-step motion and linear drift motion. Scanner drifts fast without stop in imaging mode, while it moves step by step in spectrograph mode. Tests indicated that moving accuracy is less 1 µm when the scanner moves 16µm at each step (Ahn, 2009). 34

41 Figure 2.5: The slit Slit The slit was designed and made by KASI. The initial design of the slit width was 16µm that corresponds to 0.08, necessary to achieve the intended spatial resolution of Since the amount of light was later found to be insufficient, the slit width was changed from 16µm to 32µm. The slit is now attached on the vertical table with its plane facing up, so that dusts on slit surface are one of the most serious problem. In order to block dusts, a slit cap was installed in addition. Figure 2.5 presents the slit that we used Collimating/Imaging Mirror FISS uses an off-axis mirror which originates from a portion of paraboloid mirror. The mirror makes the spectrograph configuration simple, and we don t need to concern about chromatic aberration which occurrs in a lens-based optics system. The mirror mount can control its position in 3 axes. The mirror is made of zerodur and is coated by enhanced aluminum. It has a radius of 300mm. Figure 2.6 presents the collimating/imaging mirror. 35

42 Figure 2.6: A collimator/imaging mirror Disperser The disperser is the most important component in the spectrograph. Usually prism or diffraction grating is used as a disperser. The performance of diffraction grating depends on groove density and blaze angle. Echelle grating is a kind of diffraction grating. It is different from normal grating in many ways. Echelle has low groove density (below 316 grooves/mm), high blaze angle, high efficiency and wide spectral range. To achieve wide spectral range, a prism is sometimes used as a cross disperser, but it is not used in our instrument. The Echelle grating we used is R2 (the tangent of the groove angle, blaze angle of 63.4 ), groove density of 79 grooves/mm and made by Richardson Lab. The grating surface is coated by aluminum and is overcoated by a thin layer of magnesium fluoride (MgF2) to prevent the oxidation of the aluminum. We only use spectral orders of 34th for Hα, 26th for CaII8542, respectively. The disperser is mounted on a rotating motor which is controlled by a computer. It enables the grating to fix at or rotate to the specific angle during observation. The resolution of grating angle is 1.98 /step. This value corresponds to about 1 pixel (16µm) on CCD chip. 36

43 Figure 2.7: An Echelle grating Filter An important limitation of the Echelle grating is order-overlap at high spectral orders. To avoid this order-overlap, we use bandpass filters. Free Spectral Range (hereafter, FSR) is defined as a criterion of order superposition. It is the largest wavelength range not to overlap with the spectral range in an adjacent order. If the (m + 1)th order of the wavelength (λ) and mth order of λ + λ lie at same angle, FSR is as follows: F SR = λ m. (2.14) For example, the FSR is 193Å at Hα (λ=6562.8å) and m=34. The higher the spectral order is, the lower FSR is. If we use the filter whose bandpass is smaller than FSR, we can observe the specific order that we want. For Hα, we used bandpass filter whose FWHM is 100Å. For CaII8542, FSR is 328Å so that the filter with the FWHM of 250Å is used. The transmission curves of the filters are depicted in Figure 2.8. Dashed lines represent the position at each wavelength. Filter transmissions at Hα and CaII8542 are 0.66, 0.65, respectively. 37

44 Figure 2.8: Filter transmission of Hα (left) and CaII (right). Dashed lines represent the wavelength at Hα and CaII8542, respectively. Figure 2.9: A filter adapter The position where the filter is installed is quite important. To avoid blocking light rays or internal light scattering, we put the filter at the front of the CCD. It is attached to the camera with a ring adapter. To prevent the fringe pattern which may be caused by interference between entrance light and reflection light, we have slightly tilted the filter on the mount. Figure 2.9 shows the filter adapter including a filter inside. 38

45 Table 2.6: Specifications of CCDs DV887 DV885 CCD Type back-illuminated front-illuminated active pixels pixel size (µm) 16 8 image area (mm) max readout rate (MHz) EM Gain dynamic range (Bit) CCD Cameras There are many kind of CCD cameras for astronomical observation. However, it is hard to find the cameras for specific purpose like solar observation. To catch images of faint and fine structures on the Sun rapidly, a camera should have high time cadence, high quantum efficiency and low readout noise. To satisfy these criterion, we selected DV887 for Hα and DV885 CCD for CaII8542 by Andor Company (see Figure 2.10). Main specifications are described in Table 2.6, and quantum efficiencies are depicted in Figure Two cameras are electron multiplying CCDs (EMCCD) in which a gain register is placed between the shift register and the output amplifier. The gain register consists of lots of stages, and the electrons are multiplied in each stage, so that it can get thousands of electrons from one electron with low readout noise. 2.7 Integration Integration is the process to collect each component of FISS. Integration process contains both connecting hardware components into one and merging each program to control hardware into one. Two processs are described here. 39

46 Figure 2.10: DV887/DV885 EMCCD Hardware The operation of FISS is controlled by a computer which connects the scanner, the slit cover, the grating, the CCD mounts and the CCD cameras. Figure 2.12 shows the diagram of the hardware connection in FISS. Basically a PC controls the FISS system using onboard PCI cards. In case of the grating, CCD mounts and scanner control, a servo drive is used to deliver their command from the PC to the motor. The control boxes for scanner and grating motor have been constructed in KASI. Each control box includes a power connector and a servo driver to deliver the command from the computer to each motor. Briefly speaking, the computer has four PCI cards which are connected to seven parts of FISS. Electric power is supplied not only to the control boxes but also to the CCDs separately Software We have developed the control program for each device using the software development kit (SDK). Each control program is merged into one. After in- 40

47 Figure 2.11: Quantum efficiency curves of DV887 (upper) and DV885 (lower) 41

48 Figure 2.12: A layout between hardware connections. means a motor to move a CCD mount. In this picture, M 42

49 Figure 2.13: A control program of FISS 43

50 tegration of the programs, the program for the observation was developed. Labwindows/CVI made by National Instruments is used as the software development tool. Since CVI contains ansi-c basement and graphic user interface (GUI) is easier than in MS Visual C++, we could develop the program efficiently at short time. A program technique commonly used in the FISS control programs is thread. Thread is the smallest unit of processing that can be scheduled by an operating system like MS windows. It generally results from a divergence of a computer program into two or more concurrently running tasks. We used the thread function for the monitoring of CCD temperature, realtime image display, and FISS status display. The FISS program is composed of data view and several controls of the CCD, scanner, and grating/ccd motor. CCD control The CCD program consists of 2 parts: setting the CCD parameters and saving data. Parameters of a CCD camera include exposure time, shutter on/off, cooling on/off, binning and so on. All the parameters can be controlled mostly by clicking the corresponding button on the GUI panel. FISS data are saved into files as FITs format. The program saves data in 2D form for single spectrogram mode and in 3D form for imaging spectroscopy mode or flat-fileding mode. We used the CFITSIO library from NASA (HEASARC, 2009). Each data file is named by the universal time of observation. CCD parameters and other information are recorded in the FITs header of the data. After developing basic control program, we have optimized the program to deal with several problems such as processing speed, memory management, and controlling of the 2 CCDs. Memory management is crucial in the operation of FISS. If FISS take a single spectrogram, it needs only a memory size of about 512KB. However we plan to take several hundred frames per one scan, and the FISS program requires additional memory for data manipulation for the interactive observation, so that we need much bigger memory. The maximum size of memory allocable to a 32-bit operating system is about 2.1Gbyte. But 44

51 we have to keep in mind that the memory that can be used by the FISS program is much smaller than this because the operating system and other programs installed on computer occupy a good portion of the total memory. Controlling 2 CCDs may be another problem. The two CCDs that we use are from the same manufacturer, and hence we can use the same SDK for both the cameras (ANDOR, 2009). But there are many differences in pixel number, size, gain range, affecting the size of memory, data size, window size of viewer, binning rate and so on. It is also important to have 2 CCDs interlocked. Scanner control The scanner control program has been developed by using the PMAC driver (DELTA TAU, 2003). The program has four main functions: power on/off, making script for command execution, monitoring current status, and command input. When users input a few parameters on the GUI panel, a script which includes several commands to move the scanner is created. After that, commands are sent from the script file to the ROM of PMAC board, and then the scanner starts to move. They can select one of two kinds of motions: stepwise motion and continuous motion. In both the cases, the CCD and the scanner operate at the same time. If the scanner motion is stepwise, the scanner moves a specific distance and waits until data are taken by the CCD. In the continuous mode, the scanner moves without stopping, while the CCD operates regularly. The stepwise and continuous runs of the scanner correspond to the spectrograph mode and the imaging mode of the FISS. Grating/CCD motor control Grating and CCD benches are moving parts, too, hence motor control is necessary. When FISS is turned on, grating should be positioned at specific angle precisely. Furthermore, controlling the distance between imager mirror and CCD is essential to focus on clear object. Then the program of Grating/CCD control has been developed by using SDK (PARKER, 1998). The Grating/CCD control box connects four motors, grating motor, CCD focus motors, and slit cap. The program is controlled by a computer using ethernet port. The program mainly consists of 45

52 five functions, returning home position angle, moving specific angle, checking the current status, and matching between wavelength and grating angle, and command input. Viewer Viewer program consists of 3 contents: display window to view raw data from CCD or FITs file, scaling bar to see the image in detail, and display window to view current status of system and setting value, respectively. Data viewer can display raw data from CCD and FITs file both 2D and 3D. In case of 3D data, user can select frame number to see and the direction of either raster image in spectral domain or spectrogram in spatial domain. During observation, raster image at wavelength center can be displayed on viewer panel sequentially. After obtaining, we can see either Hα and CaII image by clicking radio buttons on the panel. By using scaling bar, user can see the image in detail. If the data is 3D, frame slide bar is activated, hence we can see spectrogram or raster scan image. As long as the data is loaded in the display viewer panel, intensity histogram appears in graph panel automatically at the same time. As a mouse curser moves on data viewer panel, horizontal line profile and the mouse cursor position are displayed at the histogram panel simultaneously. Of course, users can see histogram when they click the mouse on graph panel. Parameter viewer display the data number (DN), X and Y position, average, and standard deviation of the DN. If we read the data from FITs file, it shows header information written in FITs header. Also parameter viewer can show setting parameters when we set to FISS. Important parameters are displayed such as FISS mode, scan step size, time per a scan, exposure time of the CCD and so on. The status of FISS is checked in lamp color and as scripts on panel. Integration with each program All programs to control the individual components have been integrated. Each program controls its instrument, but integration is not sufficient for observation program. The program should take 46

53 into account not only mechanical controlling but also observation condition. We added several functions into the program to apply for observation. To take calibration data for bias and dark subtraction, 100 frames are acquired for each camera with the same exposure time, binning rate, cooling and so on as real observations, but with the shutter open. Frames are averaged and then only the average one is stored in hard disk on computer. Bias/dark is obtained whenever CCD setting is changed. To get flat/calibration data, the program controls the grating, CCD and scanner successively or simultaneously. Flat/calibration data are usually taken before and after observation. When flat fielding mode is chosen for observation, grating angle is fixed specific angle which locates Hα or CaII8542 line at specific positions on the corresponding CCDs. A quiet region is scanned so that about 400 frames are obtained. At that position, the averaged one is stored into the computer. The grating angle changes 7 times and the same process as above is repeated seven times. As a result, 7-averaged frames are stored in hard disk. During solar observation, observer may want to record observation condition like position, target and so on. As an observation log, input panel was made in integration program. The observation program has the capability to record observer name, the position what they want to see in heliocentric coordinate, observation target, active region number on input panel. These information are recorded in FITs header. Figure 2.13 shows the integration program to control FISS. Main program hide a few subpanel associate with setting each component. 2.8 Installation After completing integration process, we integrated all components on optical table for tests. In the beginning of installation, each component was installed on horizontal table in KASI. To attach instruments, the position information about each device is necessary. Based on design coordinate from ZEMAX, install position on optical table was decided so that we installed each instru- 47

54 Figure 2.14: The position of FISS components ment. Figure 2.14 is the position which is supplied by ZEMAX. According to drawing, the base point is the slit position. After attachment of the slit, other instruments are attached successively. The laser is used as the light source during installation. HeNe laser (6328Å) is very useful to attach and to align instrument precisely. Furthermore, the wavelength of the laser is close to Hα. After completing to install on horizontal table, all components are also installed on vertical table. During installing on vertical table, we found that weight balance of grating is not suitable and dusts between slit blades affect spectrum quality. Then balance weight on grating are attached in addition, and slit cap was also installed to protect dust. After testing in KASI, we transported FISS to BBSO and installed on vertical table in the Coudé room. Meanwhile, cable connection among the device on the optical table, control boxes, and a computer have been performed. The control computer, grating control box, and scanner control box were located in observation room near Coudé room. Finally shield case was attached to the FISS to protect dust and stray light. Figure 2.15 and Figure 2.16 show FISS on vertical table in the Coudé lab. 48

55 Figure 2.15: Installation of FISS on vertical table in BBSO 49

56 Figure 2.16: AO (left) and FISS (right) in Coudé room. 50

57 Chapter 3 Test After the installation, the instrument was tested mainly at KASI. In order to supply light to FISS, we set up the feed optics system in front of the spectrograph. At the beginning of test, as we mentioned above, a HeNe (6328Å) laser was used as the light source. After the laser test, FISS was tested by using the sunlight. Through testing, we determined the performance of our spectrograph. Here we describe the process and results. 3.1 Configuration of feed optics A light source is needed for testing FISS. The laser and the Sun are used as the two light sources of the spectrograph. These lights are parallel, so the imager which converges the lights is essential. The fed light should be F/26 light with the same F-number as FISS in BBSO. For this, an iris is needed to adjust the effective aperture of the lens and hence the F-ratio of the converging beam. A convex lens of 1 meter focal length was used as the imager. The F-ratio of the imager was adjusted by the iris attached to it (Nah et al., 2009). The advantage of using the same focal ratio is in the flexibility that lights gathered by telescopes of different focal lengths and different field of views can be fed into the same instrument without changing the optical configuration, only if 51

58 they have the same F-ratio. When the laser is used as the light source, we need a collimator to convert the light from the point-like source into a bundle of expanded parallel beam. A small refraction telescope is used as the collimator. Generally a telescope like this is used as the light collector. The bundle of light from a distant source is converged by the object glass and paralleled by the eyepiece to be fed into an observer s eye. But we used the telescope in the opposite way by reversing the direction of light: we put the laser at the original position of the observer s eye. The light from the laser then passes through the eyepiece, converges into the focus, then passes through the objective lens, and finally becomes an expanded collimated beam. Figure 3.1 is the feed optics consisting of laser, telescope, iris and convex lens. Contrary to the laser, the Sun is not stationary on the plane of sky. We tracked the Sun and feed the sunlight into the laboratory using a coelostat. The coelostat consists of two rotating mirrors that can always reflect sunlight along the specified path. With this coelostat, the solar image neither move, nor rotate even when the Sun moves across the sky. In other words the solar light is fed into the optical table in a constant direction. Figure 3.2 is the primary mirror of the coelostat we used. 3.2 Laser Test There are many reasons why we use a laser as the light source. Since the laser is bright enough and goes straight with little dispersion on the light path, we could align all instruments with same height on optical table quickly. Moreover it is a monochromatic point source; the spectrogram of a laser source taken by CCD chip is a spot both in the spectral domain and in the spatial domain. This spot helps us not only in precisely setting the distance between imaging mirror and each CCD, but also in locating the positions of the CCDs correctly. Once the alignment of laser is complete, we don t have to work to change optical configuration of FISS any longer. 52

59 Figure 3.1: Feed optics with the laser. It contains laser, telescope as the beam collimator, iris, and convex lens. 53

60 Figure 3.2: The primary mirror of the coelostat After the alignment of all components, the laser spectrum was taken and full width half maximum (FWHM) of laser spectrum was determined. A spot size is determined by 2.2 pixel. This value is satisfactory to us in that the optimal detection sampling occurs when the resolution is equal to two pixels. The corresponding spectral resolving power is about This test result shows that the instrument has a spectral resolution close to our expectation. Figure 3.3 shows the spectrogram of laser and its profile. 3.3 Sunlight Test The main objective of tests using sunlight is to check whether the solar spectra at the two wavelengths are simultaneously recorded well or not. As a matter of fact, from the tests using sunlight, we found a few problems which affect data quality. We also determined some parameters of the instrument from these tests. 54

61 Figure 3.3: Enlarged laser spectrogram (upper) and its horizontal profile (lower). In laser spectrogram, x-direction and y-direction indicate spectral domain and spectral domain, respectively. In lower plot, a dashed line presents gaussian fitting 55

62 CCD anomaly A pair of spectrograms at Hα and CaII8542 were taken simultaneously (see right image in Figure 3.6). Contrary to Hα spectrogram, the CaII spectrogram shows not only solar absorption lines but also a wavelike pattern. This fringe pattern always appears irrespective of incident angle of grating or cooling down CCD temperature in CaII8542. After some tests, we found that this strange pattern is caused by the etaloning effect that are intrinsic to all back-illuminated CCDs working at NIR and IR wavelengths. Because back-illuminated CCDs are thin devices (typically 10-20µm) which become semi-transparent in near infrared, the reflection between the nearly parallel front and back surface of these devices causes them to act as etalons. This phenomenon leads to unwanted fringes of constructive and destructive interference which artificially modulate a spectrum. At IR wavelength where a silicon is transparent enough that light can traverse the thickness of the CCD several times, the light to bounce back and forth between the two surfaces. This interference increases the effective path length in the silicon, but it also makes a standing wave pattern. At long wavelengths like 700nm above, several passes cause constructive or destructive interference. Due to slight differences in the thickness of CCDs from wafer to wafer or chip to chip fabrication, the fringe patterns are different from CCDs each other. The only way to avoid the problem is to use front-illuminated CCD hence we replaced DV887 (backilluminated) with DV885 (front-illuminated) for CaII8542. Measurement of grating efficiency During tests, spectrograms at the same wavelength were recorded at different grating angles, with different intensities. These represent spectrograms of different orders. The intensity of a spectrogram is the highest when the grating angle is closest to the blaze angle of the grating. As the grating angle deviates from the blaze angle, the intensity decreases significantly. The ratio of the intensity of the brightest spectrogram to the intensity of all the spectrograms represents grating efficiency. In a limited range of grating angle, we searched all spectrograms. After that, the spectral order of each spectrogram was identified. 56

63 The relative intensity of each spectral order can be calculated using the blaze function: with X defined by I [sinc(x)] 2 = ( sinx X )2, (3.1) X = πθ b λ (3.2) = π(sin(α ϕ) + sin(β ϕ)) b λ = mπb sin ϕ [cos ϕ d tan α+β ], 2 where b, d, α, β are the width between adjacent grids, groove density, blaze angle, incident angle, and diffraction angle, respectively. Figure 3.4 shows the plot of theoretical and measured relative intensities as functions of spectral order at Hα and CaII8542, showing rough agreements. In case of Hα, the order of peak intensity is different: 34th order is theoretically, but 35th order experimentally. From intensity distribution curve, grating efficiencies can be calculated for Hα of 34th and CaII8542 of 26th: it is 0.34 at Hα, 0.38 at CaII8542, respectively. These values affect the instrument transmission which is used for the estimation of light level (see Section 2.5). If an observer tries to observe the Sun at Hα only, it might be a good choice to select the spectral order of 35th. In case of the Echelle spectrograph installed the Vacuum Tower Telescope (VTT) at Sacramento Peak Observatory, the specification of Echelle grating is the same as FISS (79 groove/mm, blaze angle of 63.4 ), but they use the spectral order of 36th for Hα observation (Cram et al., 1981). FISS obtains two spectra at different wavelengths simultaneously, however, so that it is better to select the 34th order for Hα observation. Determination of Linear Dispersion One of the most important parameters in a spectroscopic observation is linear dispersion, the ratio of wavelength variation to the variation of distance on the detector. The linear dispersion 57

64 Figure 3.4: Efficiency curves of echelle grating at Hα (left) and CaII8542 (right) expected from Equation 2.4. In case of spectrum data, linear dispersion can be empirically determined from the positions of spectral lines on the detector. For Hα, using two Fe lines from the solar photosphere used. We found a linear dispersion of 1.2Å/mm. The wavelength coverage per pixel is 19.2mÅ, and that of the whole detector is 9.8Å. A similar process was performed for CaII spectrum and we found the linear dispersion of 1.6mÅ/mm. These values are almost the same as the theoretically expected ones. Figure 3.5 is an example of line identification for Hα. Red circles are line of interest which we want to use. Simultaneous observation with two CCDs Simultaneous recording of two spectra was tested in KASI. Figure 3.6 shows a pair of spectrum at Hα and CaII8542, respectively. CaII8542 spectrum shows a wave-like pattern which was caused by a CCD anomaly in IR. In addition to the combination Hα and CaII8542, we tested another pair: CaIIK and Hα. The light level at CaIIK, however, is so weak that low S/N data is taken. This is because the convex lens and the mirror in the feed optics have very low transmission below 4000Å. 58

65 Figure 3.5: The solar spectrogram (left) and its line profile (right) for Hα Figure 3.6: Two spectrograms of Hα (left) and CaII8542 (right) at the same time. 59

66 Interlocking between CCD and scanner Synchronization between the CCD and the scanner was tested in KASI and BBSO. The test of interlocking between a CCD and the scanner was performed by using resolution panel located at pupil in AO. If the interlocking works well, we can see a wellconstructed raster image of resolution panel illuminated by sunlight. Figure 3.7 shows such a raster scan image for the resolution panel taken at Hα center. Resolution patterns are well constructed. During the observation, there was air disturbance. On the right side of image, distortion in vertical direction appears. This fluctuation is caused by airplane on the sky at that time. We conclude that the scanner and a CCD are working so well. Figure 3.8 shows raster scan images constructed at Hα and CaII8542 at KASI on Dec 14, Exposure time, slit width, step size per scan, frame, scanning time, and FOV are 30 ms, 16 µm, 16 µm, 512 frames, 50 sec, and , respectively. We can identify several features in raster images such as faint prominence on limb, active regions, plages, and filaments. In the CaII image displays horizontal lines the CCD anomaly. From these raster images, we conclude that two CCDs and scanner are organized well enough. More interlocking tests by using FISS will be given in the next Chapter. Light level and spatial resolution A main discrepancy between the estimation by the design and the determination from observation is in the light level; the level of light detected by CCDs is too low. Moreover, current AO is not working well enough to achieve the diffraction-limited resolution. To increase the light level, we increased the slit width from 16µm to 32µm. This means that an amount of light increases twice. Secondly, 1 2 binning was applied to CCD A, and 2 4 binning to CCD B, which results in the effective pixel size of 0.16 in the slit direction. Due to the current limitation of AO capability, it is good enough to observe the Sun for a while. If the performance of AO is improved so that it can achieve the diffraction-limited resolution, the binning can be reset to the default values. Thanks to the combined effect of a wider slit and the binning, the light level is four times the original. Observing 60

67 Figure 3.7: Raster scan image of the resolution panel on May 10, On right side (square), the distortion appears in vertical direction because of the fluctuation caused by an airplane. 61

68 Figure 3.8: Raster scan images of the Sun constructed at Hα (left) and CaII8542 (right) on December 14, Horizontal lines on CaII8542 image are caused by the CCD anomaly at IR. 62

69 Table 3.1: Typical observing parameters parameters CCD A CCD B wavelength Hα Å CaII 8542Å slit width 32 µm (0.16 ) 32 µm (0.16 ) binning integration times 30ms 30-90ms scan step size 32µm (0.16 ) 32µm (0.16 ) scan frames field of view (16-64 ) 40 (16-64 ) 40 observing duration h h CCD cooling temperature 50 C 50 C FOV of flat fielding parameters are depicted in Table Conclusion After installing of all optics and moving components on a vertical table, various tests were mainly performed at KASI. A detailed position of each component was decided by laser test, and the pixel resolution was obtained. The test by sunlight using coelostat enables us to find unexpected problems and to measure spectrograph parameters of each component. Through light test, we had opportunity to improve the FISS performance. We confirmed that scanner and CCDs operate well. 63

70 64

71 Chapter 4 Early Observations 4.1 Data processing Raw data obtained with FISS are saved in files of FITs format. There are three kinds of data science, bias/dark, and flat/calibration data. Basic data processing has been performed using IDL procedures and functions. Here we describe the process. Bias/dark, flat correction The method to take bias/dark, flat/calibration file are already mentioned in Chapter 2. Science data and bias/dark data have the same bias values. Moreover, exposure time of science data is the same as that of bias/dark data. Flat/calibration data contain seven spectrograms taken at different grating angles. Each spectrogram comes out of the averaging of about 400 frames. It is necessary to separate slit pattern and nonslit pattern. Slit pattern is inferred from the average of all the spectrograms over the spectral direction. We divide seven average spectrograms by this slit pattern, and determine the non-slit pattern from them using the flat fielding method suggested by Chae (2004). Subtracting bias/dark data from science data and dividing them by the product of the slit pattern and the non-slit pattern completes basic data processing. Figure 4.1 presents examples of frames 65

72 Figure 4.1: Examples of raw spectrogram (left top), slit pattern (right top), processed spectrogram (left bottom), flat pattern (right bottom) at Hα we mentioned this paragraph. Data compression and noise filtering The data taken by FISS have usually big size. For example, if FISS sets to 400 frames per a scan, 1 1 CCD binning, data amounts to about 400MB. The longer observation runs, the bigger data it produces. Then the proper management of big size data is one of the most important things. We applied principal component analysis (PCA) for this purpose. This technique was found to be useful not only for data compression, but also for noise suppression. Briefly speaking, the a spectrum can be described as a linear combination of basic functions (principal components) so that it can be described by a set of coefficients. An encouraging property of PCA is that it is not necessary to keep all the coefficients; about 20 to 30 coefficients is good enough keeping to reproduce most features of a spectrum. By using PCA, the data size can be compressed to one 25th while reducing noise. This noise filtering is particulary useful for CaII data the S/N ratio of which is so low. Figure 4.2 shows the line profiles of raw data and compressed 66

73 data. Compressed data match raw data well except for noise. Wavelength calibration There are lots of absorption lines near the wavelengths of our interest: Hα and CaII8542. Absorption lines come from either the Sun or the Earth. If a violent activity of the Sun happens, solar absorption lines may be modified or their positions may be shifted. Even photospheric lines might be affected a little. Hence wavelength calibration against solar lines is not suitable. On the other hand, telluric lines from atmosphere of the Earth provide a good reference for wavelength calibration because atmospheric line is not affected by the Sun. By comparing spectral lines, the FISS spectra with the reference spectrum from atlas, we identified solar and telluric lines. Using the identified terrestrial lines, we determine not only linear dispersion of the spectrograph, but also the wavelengths solar lines. The values of linear dispersion per pixel and line position (Hα, CaII8542) in pixel are recorded in FITs header. 4.2 Sunspots Figure 4.3 is the raster images of a sunspot observed with FISS. The umbra, penumbra and several jets are seen clearly both at the red wing and at the blue wing. Jets are directed outward at the blue wing and inward at the red wing, respectively. Bright points are also seen around sunspot at both the wings. Figure 4.4 show the spectra containing this sunspot. There is a dark band in the horizontal direction. This represents the sunspot darker than other regions. We find that spectral lines are either much broadened or split in this dark band. This is a result of Zeeman effect on the splitting of a spectral lint into several components in the presence of magnetic field. It is well known that the line split is proportional to the strength of magnetic field and the square of wavelength(λ 2 ). The Figure shows the splitting is more pronounced in CaII8542 than in Hα. Figure 4.5 shows a group of sunspots observed on July 22, The light 67

74 Figure 4.2: The comparison of raw data profiles with compressed data profiles. In upper row, solid lines and red dashed lines represent raw data profiles and compressed data profiles at Hα and CaII8542, respectively. Differences between raw data profiles and compressed data profiles are seen in lower row 68

75 bridge is visible inside the major sunspot. During this observation, the AO was operating and short-exposure TiO images of the same region were taken simultaneously. This is the first coordinated observation of FISS and other instruments. A speckle-reconstructed TiO image of the region is shown in Figure Filaments There are two kinds of filaments; quiescent filaments and active region filaments. Quiescent filaments are stabler and bigger than active filaments. We observed both kinds of filaments using FISS. Figure 4.7 shows a portion of a quiescent filament. It shows both the spine and barb structure. The barbs protrude from the side of the filament. Figure 4.8 shows another quiescent filament. This filament appears to be patchy rather than elongated, which is unusual. Figure 4.9 is a portion of another filament. A few dark threads are seen along the body of this filament. Figure 4.10 is an example of an active region filament. Generally it is well known that active region filament are more dynamic, short lived than quiescent filament. 4.4 Active Regions Figure 4.11 presents raster images of an active region in the emerging phase. An arch filament system appears in line center images. It is a characteristic feature of a developing active region on the Sun that connects areas of opposite magnetic polarity. Individual arches have lengths of about 30,000km, height of arches of 4,000-15,000km and lifetime of each filament of 30 minutes (Bruzek, 1967). Many bright points, called ellerman bombs, are visible at the red and the blue wing images. The lifetime are less than 10 minutes. Generally its spectrum is characterized by very broad emission wings at Hα. 69

76 Figure 4.3: Scan images of a sunspot observed on June 30, FOV is Dashed lines represent a position of spectrograms in Figure

77 Figure 4.4: Sunspot spectrograms of Hα and CaII

78 Figure 4.5: Sunspot images on July 22, FOV is

79 Figure 4.6: Speckle reconstruction image of sunspot group (AR 11089) at TiO band (706nm) on July 22,

80 Figure 4.7: Scan images of quiescent filament on July 29, FOV is

81 Figure 4.8: Scan images of a quiescent filament on July 22, FOV is 75

82 Figure 4.9: Scan images of a quiescent filament on July 16, FOV is 76

83 Figure 4.10: Scan images of an active region filament on June 25, FOV is

84 Figure 4.11: Emerging flux region associate with active region on July 22, FOV is

85 4.5 Quiet Regions Figure 4.12 is a quiet region near the center of the solar disk. Dark patches and bright points appear in FOV. Although this quiet region has no conspicuous feature, dynamical features always occur in small scale. 4.6 Prominence A Prominence is quite different from its disk counterpart, a filament. Contrary to filaments, spectral lines from a prominence on solar limb appear in emission. When a prominence is the observing target, it is hard to operate AO simultaneously because there is no good reference for AO in the limb. The image motion cannot be corrected for by the AO, and as a result we obtain raster scan images that are no so regular in the slit direction. Figure 4.13 is the prominence spectra from limb inside to limb outside. 4.7 Conclusion After testing, many kinds of imaging spectroscopy data are obtained by FISS. We have observed various chromospheric features such as active regions, sunspots, filaments, and so on. The data taken by FISS are quite good for studying the dynamical characteristics of the chromosphere. Through these observations of the Sun using FISS, we have acquired many ideas for better observing the Sun and improving FISS capability. We plan to reflect these ideas to FISS as soon as possible. We believe that the data obtained by the spectrograph will help us to discover the dynamics and physical conditions of the Sun. 79

86 Figure 4.12: Scan images of quiet region near disk center of the Sun on June 25, FOV is

87 Figure 4.13: Hα (upper row) and CaII (lower row) spectrograms taken at the solar limb and a prominence outside it. 81

88 82

89 Chapter 5 Determination of physical parameters of prominences 5.1 Introduction A prominence, meaning protuberance in French, is one of most notable features seen on solar limb above at chromosphere. Prominences are classified into two types according to their morphology: quiescent and active prominences. Quiescent prominences are stabler than active prominences, and usually maintain their shapes for a long time. Their shapes are like hedgerow, curtain, arch, etc. They are of Mm length, Mm height, 4-15Mm width (Priest, 1989). Quiescent prominences are located above polarity inversion lines and consist of many fine threads. Active prominences often occur near active regions and are associated with solar flares and other violent events. They have stronger magnetic field and higher temperature than quiescent prominences. Also their shapes last for more or less short durations. Surges, sprays, postflare loops are examples of active prominences. Although quiescent promineces are globally stable, they exhibit rapid variation in their fine structure within a few minutes from ground-based and space-based observation. It has been frequently attempted to determine their physical quantities. 83

90 Temperature is one of the most important parameters. Ellison (1952) tracked the same prominence on the limb and on the disk, and obtained a kinetic temperature of 10,000-20,000 K. Jefferies & Orrall (1961) obtained a mean kinetic temperature of about 12,000 K based on Blamer and Paschen jump for quiescent prominences. Furthermore by using line width between the higher Balmer and metallic lines, they obtained the temperature of less than 8,000 K (Jefferies & Orrall, 1962). They concluded that this discrepancy is caused by temperature and density gradients inside the matter of quiescent prominences. On the other hand, Hirayama (1963) measured temperature of 4,500-8,500 K using hydrogen lines and metallic lines. Nikolsky et al. (1971) obtained T=10,000 K, ξ=6km/s using analysis of stark effect on Balmer lines (up to number 36) from quiescent prominence. Hirayama (1971) found that the prominence has temperature and non-thermal velocity gradient from central region to periphery. In central region, temperature is 6,000-7,000 K and at the outer region, it is about 12,000 K. Zhang et al. (1987) obtained T=7,500 K by applying least square method. Stellmacher et al. (2003) obtained T=8,000-9,000 K, ξ=3-8km/s from CaII8542 and HeI10830 observation. Many authors measured temperature of prominences, but the determined temperatures are somewhat different (Kawaguchi, 1966; Kubota, 1968; Morozhenko, 1974; Mouradian & Leroy, 1977; Landman et al., 1977; Mein et al., 1991; Kejun et al., 1998). The common method to determine the temperature is to measure the width of optically thin lines and to compare with each other. Note that line width contains not only the effect of temperature but also that of non-thermal velocity. If it is possible to separate between the two components well, we will be able to better estimate the prominence temperature. Our new instrument, FISS, has a capability to simultaneously take two spectra of two different lines Hα line and CaII line at 8542 Å with high spectral resolution and with high spatial resolution. Then we expect that temperature can be determined from the widths of these lines more precisely than previous researches. 84

91 In this Chapter, we estimate physical parameters at specific positions of two prominences by using a simple model. 5.2 Model of Radiative Transfer As we know well, radiative transfer equation is described by I = I 0 e τ + τ 0 Se t dt, (5.1) where τ and S are optical thickness and source function, respectively. If there is no background radiation, source function is constant along the line of sight, Equation 5.1 may be simply rewritten as I = S(1 e τ ). (5.2) Since collisional broadening is negligible in prominences, only Doppler broadening needs to be considered. Thus optical thickness has a Gaussian profile τ = τ 0 exp( λ λ c λ D ) 2, (5.3) where λ D is Doppler width and λ C is the position of line center. In the solar atmosphere, individual atoms move in random directions continuously with the average speed so it is proportional to the square root of ion/atom temperature. These random motions of ions/atmos result in line broadening called thermal Doppler broadening. The Doppler width of thermal broadening is given by λ D = λ c 2kT M, (5.4) where λ is the wavelength at rest, T is ion temperature, M is the mass of ion. In reality, line widths determined from observations are broader than thermal broadening. This is why it is believed that there are unresolved motion, called non-thermal motion. With non-thermal motion taken into account, the 85

92 Doppler width is given by λ D = λ c 2kT M + ξ2, (5.5) where ξ is the most probable speed of non-thermal motion. If we measure doppler widths of Hα and CaII lines, we can determine temperature and non-thermal velocity separately. It is enough to use doppler widths for Hα and CaII if we only want to know temperature and non-thermal turbulent velocity. We also want to obtain other parameters associated with radiative transfer equation : source function, optical thickness, wavelength offset, etc. In order to get these parameters, we apply a nonlinear least square fitting(hereafter NLSF) to the line profiles of the prominence. The free parameters, S Hα, τ 0(Hα), λ c(hα), λ D(Hα), S Ca, τ 0(Ca), λ c(ca), λ D(Ca) are obtained from the fitting by non-linear least square method. Our method is based on the assumption that Hα and at CaII8542 emission come from the same position and the same material. In Section 5.4, we will verify whether this assumption is correct or not. 5.3 Data and Analysis We observed one prominence at the east limb and another at the west limb. Prominence data have an advantage in data analysis; since there is no background radiations in prominences. On the other hand, the light from a filament on the solar disk is mixed with the background radiation, so it is hard to separate between background radiation and filament radiation. Data were taken with FISS on June 30, Figure 5.1 shows a BBSO full disk Hα image taken on the same day. Two squares refer to the fields of view we observed. The number of steps in each raster scan is 450 frames, slit width, 32µm, exposure time, 90ms for east limb, 60ms for west limb, respectively. The prominence on east limb is more quiescent and bigger. On the other hand, the prominence observed on the west limb is more active, even though smaller. 86

93 Since AO was not working well, an irregularity appeared in raster scan images across the slit direction. After selecting a few positions in the raster scan images, we analyzed Hα and CaII8542 data. Position alignment FISS uses two CCDs with different chip sizes. Furthermore, the specific FOV of each CCD is subject to optical alignment. Due to the difference in chip size and alignment, the two CCD come to have different FOVs along the slit direction. Then the co-alignment between two spectral data is essential for data analysis. The relative offset along the slit direction was found to be 7 pixels. This offset value may change whenever optical configuration changes. Aureole correction If optical configuration is ideal, there should be no light above prominences in our measurements. In reality, however, stray light is identified on the slit surface. Stray light enters the spectrograph and appears as a ghost spectrogram on CCD. Although this aureole is fainter than the prominences, it still affects the line profiles of prominences seriously. If prominence intensity is low, we cannot measure line width because of the contamination by the aureole. In order to remove stray light, we used a portion of spectrum at the location outside prominences. The selected spectrum is averaged along slit direction. This averaged spectrum is subtracted from each row of the spectrograms, which is like bias/dark correction. Figure 5.3 shows a raw spectrogram and the processed spectrogram with stray light taken out. 5.4 Results As the first step, we arbitrarily selected two positions on the east limb prominence (see Figure 5.4), an applied the NLSF to the spectra at those points. Fitting result 1 Firstly we determined a few parameters from line profiles of Hα and CaII8542 : source function (S), optical thickness (τ 0 ), wavelength 87

94 Figure 5.1: Hα full disk image on June 30, 2010 in BBSO. Two squares represent field of view of FISS. 88

95 Figure 5.2: Scan images at Hα, CaII in east (left column) and west limb prominence (right column) Figure 5.3: A raw spectrogram (upper) and a correction spectrum subtracted aureole light (lower) at Hα. 89

96 Figure 5.4: Two points at the east prominence of Hα (upper) and CaII8542 (lower) offset (λ C ) and line width ( λ D ) using NLSF. From the obtained line widths ( λ D(Hα), λ D(Ca) ), we infer the temperature and non-thermal velocity. Figure 5.5 presents line profiles in comparison with fittings for positions A and B, respectively. The Hα line profile of position A displays central reversal. Stellmacher et al. (2003) found that central reversal occurs in Hα emission line when optical thickness is greater than 4. This indicates that source function is not constant over the optical path, and self absorption exists in an optically thick layer. In spite of central reversal in position A, the fit appears fairly good not only in the wing but also in the central region. The CaII8542 profile is fit better. In comparison with Hα, the optical thickness of the CaII line is quite small. Reliability of optical thickness From fitting of Hα data, we found that optical thickness is mostly greater than 1. As mentioned before, if prominence material is optically thin, the line profile is approximated to be gaussian. Based on the optical thickness obtained from fitting, this optically thin approximation 90

97 Figure 5.5: Fitting results of Hα and CaII8542 on position A (upper row) and B (lower row). Left column represents Hα profiles and fittings, while right column shows CaII8542 profiles and fittings. 91

98 Figure 5.6: Fitting results at two positions with τ 0 -variation. Black, green, red dashed line represent a spectrum profile from data and fitting results with τ 0 of 2.5, 0.3, respectively. Blue dashed lines present τ 0 as free parameters. does not hold at Hα. We wonder whether optical thickness is determined well enough for observed line profile or not. Figure 5.6 presents plots with optical thickness. When optical thickness, τ 0, is either 0.5 or 2.5, fitting does not match observation profile in both wing and core. The model profile by NLFS setting optical thickness as a free parameter well fits observation profile not only in wing but also in core. We conclude that the optical thickness obtained by NLFS explains observation well, and here is well determined. Fitting result 2 In fitting result 1, we obtained temperature and nonthermal velocity from the line broadenings of Hα and CaII8542 that are determined from line fitting. Now we regard temperature and non-thermal velocity as two free parameters to be determined from the simultaneous fit of Hα and CaII8542 profiles. This approach holds when both the Hα line and the CaII8542 line are emitted by the same cloud of plasma with welldefined temperature and non-thermal velocity. Specifically free parameters 92

99 are S Hα, τ 0(Hα), λ C(Hα), T, ξ, S Ca, τ 0(Ca), and λ C(Ca). Figure 5.7 shows the results for fitting. Line profiles are well fit by models. If the two lines are emitted by different clouds of plasma with different temperature or different non-thermal motion, fitting may not be good enough. We confirm from the fitting that the temperature and non-thermal velocity seen at Hα and CaII8542 are the same. Fitting result 3 All of results which we showed above are based on the assumption that the light from the Sun at Hα and CaII8542 must come from the same position and the same material. Is this assumption really true? If this is true, we expect that line of sight velocity at specific position should have the same value in both Hα and CaII8542 line. In radiative transfer equation, optical thickness profile depends on wavelength offset (see Equation 5.3) that is a direct measure of line of sight velocity. We tried to obtain several free parameters including temperature, non-thermal velocity, and line of sight velocity using NLSF through the simultaneous fit of Hα and CaII8542 profiles. Free parameters are S Hα, τ 0(Hα), V los, T, ξ, S Ca, and τ 0(Ca), respectively. Results are depicted in Figure 5.8. Models well match observed profiles at each position. We conclude that the light seen at Hα and CaII8542 come from same position and material. Temperature variation with height After taking physical parameters at 2 points, we investigated the variation of temperature and non-thermal velocity with height. In Figure 5.9, a group of symbols with the same shape represents a sample. Four samples are investigated to study the variation of temperature and non-thermal velocity with height. This is a tendency of temperature increase with height (see Figure 5.10). It is well known that both temperature and non-thermal velocity increase from prominence center to its periphery (Hirayama, 1971). He found the temperature varies from 6,000-7,000K at the center, 12,000K at the boundary. In our results, however, temperatures are relatively high. The distribution of non-thermal velocity are not 93

100 Figure 5.7: Fitting results in Hα and CaII8542 at both the position A (upper) and the position B (lower). 94

101 Figure 5.8: Fitting results in Hα and CaII8542 at both the position A (upper) and the position B (lower). 95

102 Figure 5.9: Positions of interest with height in the east prominence. found in our results. Temperature non-thermal velocity distribution In order to know the range of temperature and non-thermal turbulent velocity, many positions are selected in east and west prominence. Figure 5.11 is the raster scan image in east limb prominence with positions of interest being marked. There are two kinds of plots, line width - line width plots and temperature versus non-thermal velocity plots (see Figure 5.12). Since the doppler width contains temperature and non-thermal velocity, isothermal curve and iso-nonthermal velocity curve can be plotted in the line width - line width plot. Blue lines (diagonal lines) represent iso-temperature curves and red lines (horizontal lines) indicate isononthermal velocity curve. In this plot, we found that Hα line broadening is more sensitive to temperature than non-thermal velocity, while CaII8542 line broadening is sensitive to non-thermal velocity. That is, Hα line broadening mostly reflects thermal broadening, while CaII8542 line broadening does non- 96

103 Figure 5.10: Temperature variations with height. thermal broadening. Spot symbols represent the values obtained by NLSF. Based on the iso-temperature curve and iso-nonthermal velocity curve, we found that the ranges of temperature and non-thermal velocity are 7,500-13,000 K and 5-11km/s, respectively. In the temperature - non-thermal velocity plot, an error bar is shown for each spot. The scale of error bar is too small to be seen in the line width - line width plot. The same analysis was performed in the west prominence as well (see Figure 5.13). Figure 5.14 shows the two plots. These plots show the same characteristics as the plots of the east limb prominence. Since light level is so low, it s hard to identify profiles of CaII8542 line. Furthermore, AO was not working well, so the seeing has not been compensated for effect. These conditions cause CaII line to have low S/N ratios. In the west limb prominence, there is no tendency for temperature with height. 97

104 Figure 5.11: Sampling positions of Hα (upper) and CaII (lower) in the east prominence. Figure 5.12: Doppler width ( λ D(Hα), λ D(Ca) ), temperature, and nonthermal velocity plot (left) and temperature non-thermal velocity plot (right) in the east prominence. 98

105 Figure 5.13: Sampling positions in the west prominence. Figure 5.14: Doppler width ( λ D(Hα), λ D(Ca) ), temperature, and nonthermal velocity plot (left) and temperature non-thermal velocity plot (right) in the west prominence. 99

106 Figure 5.15: The position where the line broadening is biggest at Hα (upper) and CaII8542 (lower). Determination of maximum temperature We tried to determine maximum temperature of prominence. As we mentioned above, Hα profile is more sensitive to temperature than to non-thermal velocity. Then we tried to search the position where line width is biggest (see Figure 5.15). The position is at the outer boundary of the prominence that is very faint. At this position, there is no light of CaII8542. Then determined line broadening was 0.43Å. The corresponding temperature range is 16,000-22,000 K under the assumption of non-thermal velocity of 4-11km/s. If non-thermal velocity is about 7km/s which is average speed of our result, maximum temperature of prominence we analyzed is about 20,000 K. This result is consistent with previous research (Hirayama, 1985) which was inferred from emission-line wings. 100