Final Report An investigation to determine how accurately the two broad band X-ray measurements from the NASA geostationary operational environmental satellite (GOES) can reproduce the full X-ray spectrum of the Sun Matthew Lewis Supervisor: Dr Barry Kellett 24 th March 2006
Contents 1. Introduction 1 2. Motivation 1 3. Literature Review 1 3.1. Introduction 1 3.2. Solar physics 2 3.3. GOES 3 3.4. XSM 4 3.5. X-ray detection methods 5 3.6. Solar X-ray satellites 5 4. Method 7 4.1. Acquisition of GOES and XSM data 7 4.2. Classification of XSM data and background reduction 7 4.3. Noise reduction 8 4.4. Determination of XSM and GOES response curves 9 4.5. Calibration of XSM data with GOES 10 4.6. XSM field of view 11 4.7. Comparison between XSM field of view and ratio 12 4.8. Correlations 13 5. Results 13 5.1. GOES 13 5.2. XSM 15 5.3. XSM field of view 18 6. Interpretation 21 6.1. Correlation of calibrated measurements 21 6.2. Ratio of calibrated data & XSM field of view 23 6.3. Linear Pearson correlation 25 7. Conclusions 26 7.1. Method 26 7.2. Solar Flares 27 7.3. XSM field of View 27 7.4. Correlation 27 7.5. Implication of results 28 8. Suggestions for Further Work 28 8.1. Observations for increasing solar activity 28 8.2. Different correlation methods 28 8.3 Reproduction of the solar X-ray spectrum 29 9. Acknowledgements 29 10. References 29 Word Count: 10,911 words -i-
Abstract A method is presented to compare the GOES XRS solar flux measurements in the 1-8 Å and 0.5-3 Å spectral regions to the full solar X-ray spectrum produced by XSM. The XSM spectrum has been calibrated for both the long and short wavelength XRS ion chambers so that the measurements compared do not depend upon the response of the instruments. A full X-ray spectrum has been produced, for both low and high solar activity, using XSM. During a solar flare an X-ray flux with energies above 5 kev is detected along with a strong Fe emission line at 6.7 kev. The amount of flux with energies below 5 kev also increases compared to flux levels during low solar activity. The background solar flux incident upon XSM has been measured and removed, leaving only the solar flux visible. The XSM field of view has been measured to be 94 degrees compared to stated value of 104 degrees. A calibration curve for the efficiency of XSM due to its field of view has also been produced. A Linear Pearson Correlation indicates a high correlation between the XSM and XRS instruments. The correlation is better for the long wavelength than the short wavelength flux measurements and both correlations are dependant upon solar activity, which could be due to inaccuracies in the XSM response curve at high energies. The XSM measurements are several times larger than the time equivalent XRS short wavelength measurements. As this is not true for the long wavelength measurements it would appear that the response curve for the short wavelength XRS ion chamber is inaccurate. -ii-
1. Introduction This document is the final report for an investigation to determine the accuracy to which the geostationary operational environmental satellite (GOES) can reproduce the full X-ray spectrum of the Sun. Included in the report is a literature review, which is a critical update on the original literature survey. The methods adopted to complete the investigation are fully explained and justified. Results of the investigation are shown, with detailed analysis and interpretation to their scientific meaning. Conclusions are then drawn from the results, along with recommendations for future work. 2. Motivation GOES satellites have recorded the solar X-ray flux for nearly 30 years in two broad energy ranges. If GOES data can be shown to accurately reproduce the full solar X-ray spectrum, then essentially all of the historical GOES data can do the same. This would mean that the full solar X-ray spectrum would be known for nearly all of the time that the GOES satellites have been measuring the solar flux. A full solar X-ray spectrum over this large time scale would be particularly useful, as it would allow solar scientists to compare the spectrum over several solar cycles. Motivation for a better understanding of the Sun comes from a number of sources. A major reason is that solar activity greatly effects life on Earth, in particular high-energy events such as solar flares and coronal mass ejections. These high-energy events can disrupt radio transmissions, cause power blackouts and damage satellites. By studying the Sun, the closest star to Earth, scientists can make predictions for the detailed structure and behaviour of other stars. GOES sees the whole Sun without any spatial resolution so the measurements can be directly compared to instruments used to view other stellar objects. A full X-ray spectrum of the Sun over 30 years would prove very useful when looking at other stellar objects. 3. Literature Review 3.1 Introduction The literature review is a critical update on the original literature survey produced on the 27 th October 2005 [1]. As well as critically analysing material already discussed the review includes discussion of new literature. The review analyses material published on the subjects of GOES, the X-ray solar monitor (XSM), past and present X-ray satellites, X-ray detection methods and solar physics. The critical aspect of the review is based upon the experience gained during the course of the investigation. 3.2 Solar Physics The solar magnetic cycle is roughly 11 years in duration, during which the magnetic polarity of the Suns magnetic field is reversed. A full cycle of 22 years, after which the magnetic configuration is restored, is known as the Hale cycle [2]. During the 11-year cycle, active regions i.e. sunspots, move from high latitudes ( 40 ) towards lower latitudes ( 10 ). -1-
The cyclic behavior of migrating sunspots can be understood in terms of a reversal of the solar magnetic field. An explanation to this was first proposed by Babcock in 1961 [3]. The explanation involved the initial poloidal field being converted into a toroidal field under the influence of differential rotation. Advances on the above paper suggest that the solar cycle is driven by the internal magnetic field in the tachocline at the bottom of the convection zone. This is a theory commonly known as the dynamo process. The paper by Charbonneau [4] explains the dynamo model in a clear and understandable way. The paper continues to critically analyse this model and concludes with a discussion of key questions relating to the dynamo model. The questions discussed here are ones that have not been fully answered and so put doubt into the readers mind as to whether the dynamo model is correct. The books read have many similar ideas regarding the activity of the Sun [2], [5], [6]. A summary of these similar ideas is given below. The summary concentrates on the ideas about the X-ray spectrum of the Sun and its activity in the form of solar flares. Hard X-ray images have led to the recognition of an impulsive phase near the start of most flares, which is not seen in the soft X-ray spectrum. Figure (6.16) [6], shows a double flare on the 5 th November 1980. The graph shows the impulsive phase of the flare in hard X- rays, this impulsive phase cannot be seen in the corresponding soft X-ray graph. It is agreed that the soft X-ray continuum of the quiet solar corona and active regions is caused, at least partly, by Bremmstrahlung (free-free transitions) of energetic electrons in the electrostatic field of protons. It is emitted by plasma at a temperature of a few million degrees Kelvin. The hard X-ray Bremsstrahlung emission of flares would require much more energetic conditions. If the electrons have an energy distribution like those in a hot gas (a Maxwellian distribution), then the temperatures must be massive, around 10 8 or 10 9 K. This has led to believe that the electrons are not thermal, but are excited by some process in the flare. Figure (6.18) [6] shows a schematic diagram of the magnetic field in a single flare loop. It shows how the acceleration of electrons and protons from the flare energy release point could produce emission of hard X-rays, gamma rays and ultra-violent lines. The soft X-ray emission spectrum consists of lines as well as the continuum caused by Bremmstrahlung emission. As the temperature of the plasma is so high, atoms are generally stripped of all their n=2 and n=3 orbiting electrons. Helium and Hydrogen exist in a completely ionized state. The remaining bound electrons are excited by collisions with the free electrons. Photons are then emitted by the downward transition of the excited but bound electrons. This produces the prominent resonance lines of Hydrogen-like and Helium-like ions of abundant elements in the soft solar X-ray spectrum. A solar flare is a sudden release of energy from the Sun. This energy appears as electromagnetic radiation over a wide range of wavelengths and as mass, particle, wave and shock wave motions. Flares occur in active regions and occur most often when the active region is in its developing stage. But flares can also occur when the region is decaying and has lost all of its sunspots. The primary energy release of a solar flare is thought to be due to the reconnection of magnetic fields. The idea is that a volume of plasma in the corona with magnetic field lines in one direction is brought into close contact with another volume with magnetic field lines in the opposite direction. This configuration can only exist if there is an associated current sheet located along the boundary of the two magnetic fields. The extremely high conductivity of the coronal gas allows the current to flow almost without resistance. But locally there is still a small amount of resistance. This small resistance causes heating and the dissipation of the current. The magnetic field associated with this current then disappears. Magnetic field -2-
reconnection is then accomplished through the small amount of resistivity that the plasma has. The energy that the current had now appears as heat energy or goes into the acceleration of particles. 3.3 GOES GOES are environmental weather satellites that continually watch over the Earth to provide early warnings against severe weather conditions. The National Aeronautics and Space Administration (NASA) project science website [7] provides excellent information regarding GOES. Information can be found on a wide range of topics including GOES history, project status, images and technical notes. Within the technical notes is a link to the GOES I-M DataBook [8]. The GOES I-M DataBook briefly explains the technical aspects of the spacecraft and instruments. The book is split into several sections making it quick and easy to locate the relevant information. One of the instruments used in this project is part of the Space Environment Monitor (SEM) subsystem. A brief introduction explains that the SEM measures the effect of the Sun on the near Earth solar-terrestrial electromagnetic environment, providing real time data to the Space Environment Services Centre (SESC). The SEM consists of four instruments; Energetic Particle Sensor (EPS), High Energy Proton and Alpha Detector (HEPAD), magnetometers and the X-ray Sensor (XRS). Some of the data used in the investigation is from the XRS instrument and therefore it is the XRS literature that shall be focused on. The DataBook reveals that the XRS is a telescope that measures real time solar X-ray flux. The solar flux is measured in the spectral range 0.5 3 Å & 1 8 Å with a threshold sensitivity of 5 x 10-9 W/m 2 & 2 x 10-8 W/m 2 respectively. Data was only used in the investigation when flux levels were above 5 x 10-9 W/m 2 & 2 x 10-7 W/m 2 for the same respective spectral regions. This is in good agreement for the 0.5-3 Å spectral region but is lower than the suggested threshold sensitivity for the 1-8 Å spectral region, indicating that the instrument is not as sensitive as the documents claims. The document states that X-rays are detected using two ion chambers, one for each spectral range, but it fails to explain how these ion chambers differ. The International Society for Optical Engineering (SPIE) proceedings series, volume 2812, entitled GOES-8 and Beyond, contains a session on Space Environment Monitoring. This session includes two papers written on the XRS [9], [10]. The paper by Hanser [9] is a very detailed paper that describes the design and calibration of the XRS. The document names the two ion chambers as chambers A and B, which have a detection range of 0.5 3 Å & 1 8 Å respectively. The X-ray detection range matches the XRS performance given in the GOES I-M DataBook. Table (1) in the paper, shows the dual ion chamber properties of XRS. Of interest is that chamber A has a fill gas of Xenon with a Beryllium window thickness of 20 mils (508 µm) and chamber B has a fill gas of Argon with a Beryllium window thickness of 2 mils (50.8 µm). This information, along with all of the other information presented in the table can be seen in table (4.2). The information presented here is needed for an accurate calibration of the instrument. A criticism is that the information is presented using unconventional units, the information would have been easier to use if it was presented in SI units. The graph in figure (3) shows the theoretical ion chamber response for varying wavelengths for both chambers. The graph displays the absorption for each chamber against wavelength. A plot of transition against energy would have been preferred. When compared to the theoretical response using the data from table (1) and current data from the National Bureau of standards for X-ray cross-sections [11] the graphs aligned well, although not -3-
exactly. It is possible that this is due improvements in particle physics over the last ten years, therefore improving the accuracy of the data. The paper by Bornmann [10] supports the XRS response produced in figure (3), [9] by showing a similar graph in figure (2). The graph also shows the XRS response from previous GOES satellites and is the only difference between the two figures. The document makes no attempt to inform the reader which response curve corresponds to which GOES satellite. A very detailed introduction into the XRS is provided in this paper. This introduction supports the information from [8] but also expands on this by giving details about the history of GOES and by explaining some simple solar physics. 3.4 XSM XSM is a scientific instrument on the European Space Agency (ESA) satellite SMART-1 (small missions for advanced research technology). XSM is technically a sub instrument of D-CIXS (demonstration of compact imaging X-ray spectrometer), therefore to gain a full understanding of XSM an appreciation for D-CIXS is needed. The paper published by Grande [12] explains that D-CIXS will provide high quality spectroscopic mapping of the moon. The paper reveals that a calibration device for D-CIXS is needed due to the variability of the solar flux in the energy range 0.1-20 kev. Results from this project have confirmed that the Sun is extremely variable in this energy range, even over time periods of a few minutes. This solar variability results in large variations in the production of fluorescent lines above 1.5 kev. Brief details about XSM are provided in section 2.3 entitled the X-ray Solar Monitor (XSM). The section is dominated with the requirements of the calibration device but not what the detector can actually achieve. XSM is discussed independently in the paper by Huovelin [13]. The paper is meant to accompany the previously discussed paper. It explains that the primary scientific goal of XSM is to provide a calibration source for D-CIXS, this is in agreement with [12]. But the paper also describes the independent science that can be achieved by XSM. During a solar flare XSM will measure a spectrum dominated by a flux with energies above 2-3 kev. The instrument range is very sensitive to solar flare activity, this is in agreement with the energy range stated in [12]. The wide energy range above 1 kev should, therefore, provide good sensitivity to distinguish flare characteristics, a claim which the results from this project support. The good sensitivity is a time resolution of 16 seconds and a spectral resolution of 40 ev. Although XSM is capable of measuring the solar flux with a time resolution of 16 seconds, a lower time resolution, approximately 2 minutes, is needed for good statistics when subtracting the background flux. The document states that XSM has a 105 degree field of view. From the results obtained during the project the field of view is lower and a more accurate value would be 94 degrees. It is clear from the results that the efficiency of the detector is dependant upon the position of Sun within its field of view. It is, therefore, disappointing that no statement of this fact is made and no calibration factor for this effect is included. A recent paper by Vaananen [14] again describes the primary purpose of XSM. By cross calibrating XSM with GOES the paper claims that the fidelity of the XSM data seems highly likely. Results from the project agree with this statement. The increase and decrease in flux measured by XSM occurs at exactly the same time as that measured by GOES. The amount that the flux varies also seems comparable. However, there are still discrepancies between the instruments for high solar activity. During high activity, the flux detected by XSM appears to be greater than the flux recorded by the GOES XRS. Several explanations can solve this -4-
discrepancy, one of which is that the theoretical response for XSM at higher wavelengths is inaccurate. 3.5 X-ray Detection Methods GOES XRS is an example of a gas ion chamber X-ray detector, making use of two different Noble gases, Argon and Xenon. The physics behind gas ion chamber X-ray detectors is described in the paper by Thompson [15]. A typical detector consists of a rectangular gas cell with thin entrance and exit windows. As discussed earlier the entrance widow for the XRS is made from Beryllium with the thickness for chambers A and B being 20 and 2 mills respectively. An electric field of approximately 100 V/cm is applied across two parallel plates inside the detector. Photoelectrons, Auger electrons, and/or fluorescence photons are produced when X-rays in the beam interact with the chamber gas. The energetic electrons produce additional electron-ion pairs by inelastic collisions and the photons either escape or are photo-electrically absorbed. The electrons and ions are collected at the plates and the current is then measured using a low-noise current amplifier. The efficiency of the detector can be calculated from the active length of the chamber, the properties of the chamber gas, and the X-ray absorption cross-section at the appropriate photon energy. Once the efficiency is known, the photon flux can be estimated from the chamber current and the average energy required to produce an electron-ion pair. The XSM detector is an example of a high purity Silicon (HPSi), PIN diode detector [13]. Again the paper by Thompson [15] describes the physical principles behind this detection method. As an X-ray photon interacts in the intrinsic region, it produces tracks of electronhole pairs. Due to the presence of the electric field the pairs quickly separate, drifting to the detector contacts. This increases the potential difference between the contacts, thus drawing a measurable current to remain in equilibrium. The book by Tipler [16], explains the background details of semiconductors. The book starts off at a very basic level by explaining key physics such as P and N type semiconductors. It then advances further by explaining how semiconductors can be used as photon detectors. Although the book gives a clear and understandable presentation in the details of semiconductors, it does not explain how they can be used as a specific X-ray detector. 3.6 Solar X-ray Satellites There have been numerous satellites with missions to study the Sun. The book by Lang [17] gives an excellent overview of the satellites that have been used to observe the Sun. This book specifically focuses on the more famous satellites over the past 10 years. These include the Solar and Heliospheric Observatory (SOHO), Ulysses and Yohkoh satellites. The first chapter contains information regarding the instruments onboard these three spacecrafts. The book simply gives the name of each instrument and a short description of the measurements it will be taking. There are no technical descriptions on any of the instruments. Of particular interest to this project is the Gamma Ray Burst (GRB) instrument on Ulysses and the Soft X-ray Telescope (SXT) in the Yohkoh payload. Comparing measurements taken by these two instruments to data acquired using XSM and the GOES XRS would be a worthwhile task as all of the instruments have a similar energy range. Unfortunately the Yohkoh satellite was lost on 14 th December 2001 so current XSM/XRS data cannot be compared. But the Ulysses data can be compared as the satellites lifetime has been extended until 2008. -5-
The Ulysses satellite is discussed independently in the book by Balogh [18]. The book focuses on the science from the mission over the last 10 years. The technical aspects of the instruments are not discussed. Of interest is the unique orbit of the Ulysses satellite. The satellite has a polar orbit of the Sun rather than the conventional equatorial orbit, adding a completely new dimension to the way the Sun can be viewed. Due to the unusual orbit of the Ulysses satellite, careful consideration would have to be taken when comparing data directly to XSM/XRS data as the instruments would not be measuring exactly the same flux. Another satellite that is currently studying the Sun is RHESSI (Reuven Ramaty high energy solar spectroscopic imager). The RHESSI website [19] provides an enormous amount of information about the satellite through well organised links. The Data link provides a page that allows the user to view papers written on the satellites instruments. The paper by Smith [20] is written on the RHESSI spectrometer. This paper is a substantial and very detailed document, which introduces the reader to the spectrometer before explaining the components that make up the instrument. The document then finishes with a large performance and calibration section. The spectrometer was designed to study high-energy emission from solar flares over the broad energy range 3 kev to 17 MeV. The range is a lot greater than that of GOES and XSM. According to the paper, a high-energy resolution is needed to make advances in the spectroscopy comparable to the advances the high-angular-resolution rotating modulation collimator (RMC) would make in imaging. The document states that the resolution would be about 1 kev for the 3-100 kev energy range. Compared to the 40 ev resolution of XSM, the RHESSI resolution is poor. The RESIK (Rentgenovsky Spektometr s Izognutymi Kristalami) Bragg crystal spectrometer is an instrument in the payload of the Earth orbiting satellite CORONAS-F. RESIK is the subject of the paper by Sylwester [21], which outlines the objectives and concepts for the instrument. RESIK is a bent crystal spectrometer covering four soft X-ray spectral ranges (3.369-3.879 Å, 3.821-4.326 Å, 4.307-4.890 Å, and 4.960-6.086 Å). Results from XSM show that flux from the solar continuum is detected mainly at wavelengths above 2 Å. Therefore, this makes the spectral range of RESIK extremely sensitive to the solar continuum. The thermal X-ray spectrum for the 23 rd April 2003 solar flare is the subject of the paper by Dennis [22]. Data used in this paper comes from the RHESSI satellite and the RESIK spectrometer. Section 2 describes the theoretical X-ray spectrum in the 3.8-10 kev energy range due to line and continuum emission from thermal solar flare plasmas. The conclusion to this chapter is that the thermal spectrum consists of free-free and free-bound continua and two line features at 6.7 and 8 kev that, for flare temperatures less than 20 MK, are made up of He-like Fe (Fe XXV) lines and Fe XXIV dielectronic satellites. Figure (4) shows the X-ray spectra with RESIK and RHESSI with the clearly defined Fe line at 6.7 kev. This result is verified by the XSM data, which also shows a strong Fe line at 6.7 kev. The near-earth asteroid rendezvous (NEAR) shoemaker spacecraft landed on the asteroid 433 Eros on 12 th February 2001. Using remote-sensing X-ray fluorescence spectroscopy, major element ratios for the surface composition of the S-class asteroid can be determined [23]. GOES XRS data has been used to determine the solar temperature, which is needed for analysis of the results. The paper claims high-spectral-resolution observations of X-ray lines of Ca and Fe by the Bragg crystal spectrometer on the YOHKOH indicate that flares often include higher temperature components than the average temperature determined by GOES. Due to these results, the paper claims that the GOES temperatures are not sufficiently accurate to quantitatively determine abundances from NEAR XRS spectra. -6-
From the results of this investigation the GOES flux measurements for harder X-rays appear to be to low during high solar activity. If a temperature was assigned to these flux values then it would also appear to be too low, as harder X-rays are the product of a higher temperature. Unlike NEAR, many scientific results have used the GOES XRS data to imply a temperature. If the GOES data is found to be inaccurate then it could mean that many of the interpretations of scientific results could be wrong. 4. Method 4.1 Acquisition of GOES and XSM Data The GOES data was acquired by reading in the appropriate sections from 12.TXT files, 1 file for every month of data, using the computer programming language IDL. XSM data was acquired in a similar way, with the exception that XSM doesn t operate continually, so the data is periodic instead of continuous. Within the GOES data are the occasional bad data points. These are filled with a flux level of 3.27 x 10 4 W/m 2 so as to be easily recognisable against the good points. Using IDL these flux levels have been removed so that only the good data is used in the investigation. 4.2 Classification of XSM Data and Background Reduction Once acquired, the XSM data was classified into solar, calibration and background spectra. The calibration spectra, due to a small 55 Fe source attached to the back surface of the shutter and covered with a 5µm thick Ti foil, can be seen every time XSM opens and closes its shutter. The solar spectrum can be seen when XSM is operating and the Sun is within its field of view. If XSM is operating but the Sun is outside its field of view then the background flux is detected. Certain energy channels in XSM have a distinct count range that varies depending on the type of spectrum it contains. Once these count ranges were known, parameters were set to enable classification of each individual spectrum. The parameters can be seen in table (4.1) below. Once the corrupt data was taken out, each spectrum in the data set was classified via a series of complicated loops and if statements that related to the energy channel parameters. Energy channel range Counts have to be greater than Spectra type 0 509 10000 Bad 40 99 1000 Solar 100 149 1000 Calibration 350 509 18 Bad Table (4.1): Parameters to classify spectra type The background flux can be interpreted as a combination of two things. One is the sky background, which affects all the data and the other is a particle background that only a few of the data sets are affected by. The particle background flux was acquired when the sky background was subtracted from the solar spectrum taken on the 27 th April 2004. The result left some background still present. By fitting lines to remove this extra background flux, the particle background flux was formed. The sky background flux was determined using data from XSM on the 8 th July 2004, when the Sun was outside the XSM field of view. -7-
The sky background is mainly due to the temperature of Universe, which produces a constant background flux. The particle background consists of cosmic rays, which are roughly constant in time and space, and a solar proton flux, which is not constant but highly variable with a strong dependence upon solar activity. An image was produced to show the number of counts every 16 seconds recorded by XSM in every channel as a function of time. Due to the low number of counts recorded by XSM it was necessary to combine adjacent time periods to achieve higher counts and good statistics for the background subtraction. For every data set, the total time that data was recorded was split up in to a number of bins. The number of bins is set so that the data has a time resolution of 2 minutes. In each time bin the counts per energy channel were totalled. A plot was then made with the energy channel along the y-axis, the time along the x-axis and the number of counts, as a function of the time and energy channel, shown as a varying colour. The background flux was taken away from the plot, leaving only the solar and calibration spectrum visible. Over the top of the image, the GOES flux for the same time period was displayed. This enables the XSM and GOES data to be directly compared. The GOES flare classifications were included on top of the plot to enable a comparison of solar activity between different data sets. 4.3 Noise Reduction After the background flux had been removed a small number of counts were still detected in the higher energy part of the spectrum. These counts, which occurred in random time and energy bins, were due to the failure of the background reduction method to remove all of the background flux. Therefore these counts still had to be removed. Although the removal of this noise was important it was equally important not to confuse the noise with actual solar measurements. A programme was created in IDL to remove this noise. The programme set the number of counts in a certain time/energy bin to zero if all of the adjacent time/energy bins had zero counts. Figure (4.1) is a visual representation of how the programme classified noise and how it was removed. Energy Channel Time (a) (b) Figure (4.1): Process of noise reduction. The dark square in (a) represents a data bin that has a certain number of counts, the white data bin surrounding it have no counts in them. The program created in IDL will recognise the black data bin as noise and so will reset the data bin to zero counts, as shown in (b). -8-
4.4 Determination of XSM & GOES Response Curve The GOES XRS response curves, shown as figure (A.1) in appendix A, have been digitised so that they can be used in the calibration of the XSM data. The figure, which contains the response curves for both ion chamber detectors, was scanned into the computer and displayed as a.jpeg picture. Using a program created in IDL the two calibration curves were digitised using approximately 100 points for each curve. By assigning the correct values to the axis of the graph and fitting a polynomial to both sets of data points, an accurate digitised version of both calibration curves was formed. The energy response of each ion chamber is produced using a Beryllium window and a fill gas. Chamber A (0.5-3 Å) has a Beryllium window thickness of 508 µm compared to the 50.8 µm thickness of chamber B (1-8 Å). The window thickness provides the chamber response for the lower energy X-rays. A large window thickness will provide a barrier against lower energy X-rays, thus the 508 µm window of chamber A has a cut off wavelength at 3 Å compared to the 8 Å cut off wavelength for the 50.8 µm window in chamber B. The response of each ion chamber at higher energy X-rays is due to the type of fill gas used, its pressure inside the chamber, and the field of view gas thickness. Xenon at a pressure of 24 kpa, thickness of 5.80 mg/cm 2 and with a ionisation energy of 22.0 ev produces the 0.5 Å cut off wavelength of chamber A. Argon at a pressure of 106.7 kpa, thickness 6.831 mg/cm 2 and with a ionisation energy of 26.2 ev is responsible for the 1 Å cut off wavelength of chamber B. Within the two calibration curves are several distinct features that interrupt the smooth curves. These are absorption edges and are characteristics of the fill gas within the chamber. Xenon has two absorption edges within the energy response of the detector. The edge at 4.78 kev is due to L(Ш) absorption edge and the edge at 34.56 kev is due to the K absorption edge. Argon only has one absorption edge, visible at 3.20 kev and is a K absorption edge. The K absorption edge occurs when an electron is ionised from the K shell or n=1 quantum number. This is the innermost electron shell and therefore has the greatest ionisation energy. The L absorption edge occurs when an electron is ionised from the L shell or n=2 quantum number. Due to fine structure splitting the L absorption edge occurs three times, producing L(I), L(II) and a L(III) absorption edges, which all occur at slightly different energies. Fine structure splitting occurs due to multiple electrons being present in an atom, therefore the quantum numbers due to orbital angular momentum, l and total angular momentum, j must also be considered. Table (4.2) shows the quantum number configuration for both the K and L energy levels. Energy Level n l j K 1 0 1/2 L(I) 2 0 1/2 L(II) 2 1 1/2 L(III) 2 1 3/2 Table (4.2): Quantum number configuration for various energy levels in an atom A theoretical version of the two calibration curves was produced to verify that the digitised version of the calibration curve was accurate. This theoretical version, again produced in IDL, was formed using the ion chamber properties from table (4.3). For each ion chamber two plots were produced, one for the effect of the transition due to the Beryllium window and one for the fill gas. The two effects were then combined to produce a final plot showing the ion chamber response of the detector. Within the theoretical model the -9-
Beryllium window thickness and fill gas parameters of the ion chamber detectors could be easily changed to vary the effective area response at different wavelengths. Item A Chamber B Chamber Nominal X-ray detection range 0.5 3 Å 1 8 Å Beryllium window thickness 508 µm 50.8 µm Fill gas and pressure Xenon/24 KPa Argon/106.7 KPa Window area 5.80 cm 2 1.90 cm 2 FOV direction gas thickness 5.051 mg/cm 2 6.831 mg/cm 2 Energy to produce electron-ion pair 22.0 ev 26.2 ev Table (4.3): GOES XRS gas ion chamber properties The predicted spectral response of the XSM detector, shown in appendix C, figure (C.1), was scanned into the computer and displayed as a.jpeg picture. The response curve was then digitised using the same method used to produce the XRS response curves. A theoretical model, using parameters from table (B.1) in appendix B, was again used to verify the digitised response curve. The K absorption edge of Silicon, which occurs at 1.84 kev, can be seen in the response curve. 4.5 Calibration of XSM data with GOES In order to directly compare the XSM and GOES flux, the XSM measurements had to be calibrated for both instruments. XSM detects the photons from the Sun and displays the results as counts. So, in order to calculate the solar flux seen by XSM, the counts produced by XSM were put through the XSM response curve to display the photon flux that was incident upon XSM. Both XSM and GOES look at the Sun from approximately the same distance and position. Therefore, if the XSM calibration is correct, the calculated flux incident upon XSM should be the same as the flux incident upon GOES. The flux from the Sun is detected by the GOES XRS in two ion chamber detectors. Each detector converts the photons into a single count number, so as to produce one number for each spectral range. This process was recreated by passing the calculated flux incident upon XSM/GOES through the XRS response curves. This process takes into consideration the response of the XRS dual ion chambers allowing direct comparisons between XSM and GOES to be made. A visual representation of this process is shown in figure (4.2). Figure (4.2): Diagram showing a visual representation of the calibration process so as to allow direct comparisons between the XSM and GOES data The GOES data used in this investigation has been converted back into a flux, therefore the same must be done for the XSM data calibrated for XRS. The flux can be calculated -10-
using equation (4.1) for the short wavelength and equation (4.2) for the long wavelength spectral regions. Flux Flux F C XSM WE XSM = H G I A E t K J A F C XSM WE XSM = H G I A E t K J B A B (4.1) (4.2) E XSM is the mean energy of each XSM energy channel, C XSM is the number of counts in the individual time/energy bin, W is the width of the XSM energy channel, A is the window area and t is the width of the time bin. The energy to create an electron-ion pair also has to be considered as the XRS records a count for every electron-ion pair created in the fill gas. The energy is a characteristic of the fill gas, where the energy to produce an electron-ion pair for Xenon, E A, is 22.0 ev and for Argon, E B, the energy is 26.2 ev. GOES flux measurements give a single flux number for the spectral ranges 0.5-3 Å and 1-8 Å. The XSM data, although calibrated for the GOES ion chamber detectors, still stores the flux in 512 energy channels. To extract a single flux measurement for each time bin, the flux in every energy channel was summed together. This process was completed for both spectral regions, thus giving two broad band measurements for the XSM data. 4.6 XSM field of View Solar flux measurements from XSM are dependent on where the Sun lies with respect to the detectors field of view. The effective area of the detector decreases as the angle between itself and the Sun increases, hence XSM is most efficient when it views the Sun directly. SMART-1 has the ability to be rotated in 3-dimesions, i.e. along the X, Y and Z axis, but fortunately the Y angle remains constant at 90 degrees therefore reducing the rotation to 2- dimensions. As the position of SMART-1 is constantly recorded, the position of the Sun relative to XSM can be calculated. Figure (4.3) shows the co-ordinate system that has been assigned to SMART-1, with the +X direction pointing to the Sun. XSM is mounted onto the +X panel of SMART-1 and is tilted at 45 degrees with respect to the X-axis. -11-
Figure (4.3): Co-ordinate system for SMART-1. XSM is positioned on the Sun facing +X panel of SMART-1. XSM has a circular field of view, so any rotation about the X-axis only acts to rotate the detector. This rotation only affects which side of the detector the Sun is positioned on and not the angle itself. Therefore, the only rotation that changes the XSM field of view is about the Y-axis, which results in a change to the X co-ordinate. XSM has been calibrated for its field of view and the resulting calibration diagram can be seen in figure (D.1) in appendix D. As only a change in the X co-ordinate will result in a change in the XSM field of view, only the horizontal axis of the calibration diagram needs to be considered. The horizontal axis of the calibration plot was digitized in the same way as the GOES ion chamber response plots. Two calibration curves were formed by fitting a polynomial to the digitised calibration data, one curve was for the right hand side and the other for the left hand side of the XSM field of view. 4.7 Comparison between XSM Field of View and Ratio For every XSM data set the GOES calibrated XSM flux was divided by the time equivalent GOES true number. This produced two sets of ratios, one for the short and one for the long wavelength measurements, at a time resolution of 2 minutes. The XSM field of view was also averaged to a time resolution of 2 minutes so that comparisons could be made between the ratio and the XSM field of view. In order to identify the activity of the Sun for each particular ratio, the GOES flare classifications were introduced, table (4.4). Every ratio number was marked with the flare classification using the long wavelength GOES flux level for that particular time. Two plots were then produced, one for each spectral range, showing the ratio against XSM field of view. Different flare classifications were represented using different colours, and different symbols were used depending upon which side of the XSM field of the view the Sun was positioned. The two calibration curves for the XSM field of view were then placed over the top of the graph so a comparison of the calibration to actual data could be seen. GOES XL Flux level (Wm -2 ) 1x10-6 F < 5x10-6 5x10-7 F < 1x10-6 1x10-7 F < 5x10-7 F < 1x10-7 Flare Classification C1 C4 B5 B9 B1 B4 A Table (4.4): GOES flare classifications -12-
4.8 Correlations A linear Pearson correlation was used to produce quantitative results for the correlation between the GOES and XSM data. The linear Pearson correlation coefficient R, is a scalar quantity between the values of -1.0 & 1.0 and is defined as the ratio of the covariance of the sample populations, X & Y, to the product of their standard deviations, equation (4.3). (4.3) The sample populations used in the correlation were the XSM data calibrated for the XRS and the GOES true numbers for the two different spectral regions. A linear Pearson correlation coefficient was produced for all of the XSM data sets for every flare classification. A benefit of this was the effect that solar activity has upon the correlation could be measured. 5. Results 5.1 GOES The solar flux recorded by the GOES satellite between January and December 2004 can be seen in figure (5.1). The blue line shows the solar flux in the 0.5 3 Ả spectral region and the red line shows the recorded solar flux in the 1 8 Ả spectral region. Figure (5.2) shows the ratio between these two flux lines for the same time period. As expected the ratio increases as solar activity increases, but the ratio is also high when the solar activity is low. When the solar flux is low the instrument is not sensitive enough to record accurate results. Therefore the data has been removed when the flux levels fall below 5x10-9 WM -2 and 2x10-7 WM -2 for the spectral regions 0.5 3 Ả and 1 8 Ả respectively. Figure (5.1): Solar flux between January and December 2004-13-
Figure (5.2): Ratio of short wavelength to long wavelength flux measurements between January and December 2004 Figure (5.3) shows the data that remained once the low flux levels were removed. Figure (5.4) shows the new ratio between the two flux levels. The ratio is now only high when the solar activity is high. The ratio peaks at a value a value of 0.43, this is when GOES recorded the highest solar flux levels during day 194. Figure (5.3): Solar flux between January and December 2004, with low flux levels removed Figure (5.4): Ratio of short wavelength to long wavelength flux measurements with low flux levels removed Figure (5.5) displays the change in solar flux between day 95 and 105 in the year 2004, with a solar flare highlighted in yellow. Figure (5.6) shows how the flux ratio of GOES two flux measurements changes with varying solar activity over the same time period as the graph in figure (5.5). The red dots show the correlation during the highlighted solar flare. -14-
Figure (5.5): Solar flux between day 95 & 105 with a solar flare highlighted in yellow Figure (5.6): Ratio of short wavelength to long wavelength flux measurements against solar activity with a solar flare highlighted in red. The spacing of the dots in figure (5.6) show that the X-ray solar flux rises sharply during the flare before returning to normal levels at a slower rate. During a solar flare the highest ratio values occur at the beginning of a flare. This indicates that the short wavelength flux increases faster than the long wavelength flux. The ratio is always greater during a solar flare than during normal solar activity, which implies that the short wavelength flux must increase by a greater amount than the long wavelength flux. When the highest long wavelength flux measurements were recorded, the ratio value had already started to decrease. This indicates that the difference between the flux levels for the two spectral regions had started to return to a normal separation. At a long wavelength flux levels below 3x10-7 Wm -2 the dots form very distinct, solid curves. This indicates a quantisation of data and occurs when the solar flux levels approach the threshold sensitivity of the detectors. Measurements for 1-8 Å spectral region, taken where the recorded flux levels go below 3x10-7 Wm -2, may not be accurate. 5.2 XSM Figures (5.7) and (5.8) show the solar spectrum detected by XSM when the solar activity was low and high respectively. In both figures the purple lines show the solar spectrum detected by XSM, the black lines show the total background spectrum and the red lines show the background corrected solar spectrum. The other lines in the figure (5.8) are evidence of the background fitting to match the solar spectrum at the higher energy values. Here total background spectrum is now composed of sky (green line) and particle (yellow line) -15-
components. When there are no green or yellow lines, as in figure (5.7), the total background spectrum is just equal to the sky background. The letters in the far right corner of the figures are values from the solar and background spectrum. The meaning of these numbers can be seen in table (5.1). Figure (5.7): Solar flux incident on XSM for normal solar continuum. The purple line is the recorded spectrum, the black line is the background spectrum and the red line is the background corrected solar spectrum. The difference in the solar spectrum for varying solar activity is very noticeable in figures (5.7) and (5.8). During a solar flare photons with energies above 5 kev are detected and a distinct emission line for Fe can be seen at 6.7 kev. Also, the number of photons detected with energies below 5 kev increases compared to the number detected during the normal solar continuum. -16-
Figure (5.8): Solar flux incident on XSM for higher solar activity. The green, black and yellow lines are due to the background fitting. Letter Meaning A Average number of counts recorded in sun spectra every 16 seconds 300-500 B Average number of counts recorded in background spectra every 16 seconds 300-500 C Average number of counts when background is subtracted from solar spectra every 16 seconds 300-500 D Average number of counts recorded in sky background spectra every 16 seconds 300-500 SC Sky background coefficient N/A PC Particle background coefficient N/A Table (5.1): Explanation of letters used in figures (5.7) and (5.8) Energy channel Figure (5.9) shows the counts per second recorded by XSM in every channel as a function of time. The bold red line is the long wavelength GOES flux for exactly the same time period. The straight red lines show the GOES flare classifications, the values for which are shown on the right hand side of the graph. The figure consists of two graphs, the bottom one is an exact replica of the top one but with the exception that the background flux has been removed. The background flux, recorded by XSM on the 8 th July 2004, is displayed in figure (5.10). The first data that can be seen is the calibration data, obtained from the opening of the shutter door. The solar flux can then be seen, after which Sun moves out of the XSM field of view leaving only the background flux visible. -17-
Figure (5.9): Image produced by XSM data showing a B5 flare. The bottom graph is a duplicate of the top graph but with the background flux subtracted. Figure (5.10): XSM image showing the background, solar and calibration flux 5.3 XSM Field of View The effect that the Suns position, with respect to the XSM field of view, has upon the number of counts produced by XSM can be seen figure (5.11). The figure consists of four graphs, the top graph shows the number of counts produced by XSM against time, the middle three graphs show the X, Y & Z angles of SMART-1 and the bottom graph shows the position of the Sun within the XSM field of view. As the position of the Sun moves towards the centre of the field of view the number of counts produced by XSM increases. -18-
Figure (5.11): Number of counts produced by XSM and the position of Sun with respect to the XSM field of view The XSM field of view is different depending on which side of the detector the Sun is positioned. Figure (5.12) shows data taken when the Sun was positioned on the right side of the detector. The dotted line in the figure shows that the exact angle when XSM starts to record solar flux is 42 degrees. A similar plot is shown in figure (5.13) but this time the Sun is positioned on the left side of the detector. On this side the angle at which XSM starts to record solar flux is 52 degrees. -19-
Figure (5.12): Solar spectrum detected when Sun reaches 42 degrees on the right hand side of the XSM field of view Figure (5.13): Solar spectrum detected when Sun reaches 52 degrees on the left hand side of the XSM field of view -20-
6. Interpretation 6.1 Correlation of Calibrated Measurements Once XSM data has been calibrated for the GOES XRS, comparisons between measurements from both instruments can be made. Figure (6.1) shows both GOES calibrated XSM data and GOES true numbers for both spectral regions. The red and blue solid lines are the GOES long and short wavelength flux levels respectively. These flux levels have been averaged over a time period of two minutes to match the time resolution of the XSM data. The purple and green dots represent the GOES calibrated XSM data, again for the short and long wavelengths. Due to XSM being switched off there are no measurements in the middle and at the end of the graph. The calibration data, which would normally be present, has been removed. Figure (6.1): Averaged GOES flux for short and long wavelength flux measurements compared to GOES calibrated XSM data The XSM data matches very closely for long wavelengths, but for the short wavelength measurements the flux is several times greater than the GOES flux. During the time that the XSM data was recorded the Sun was at approximately 2 degrees from the centre of the XSM field of view, therefore the data needs very little correction for the field of view efficiency. Even though the XSM data matches very closely for the long wavelength measurements, it still appears to be higher than GOES flux during the highest solar activity. The difference in the short wavelength measurements is also greater during higher solar activity. In order to determine a correlation between the XSM and GOES measurements, a ratio of GOES calibrated XSM data to GOES data was plotted against time for both the spectral regions. The ratios, for the same data used to produce figure (6.1), are displayed in figures (6.2) and (6.3). The red dots represent the ratio for when a flux level greater than or equal to a C1 flare was recorded by GOES, the blue dots represent a flux level lower than a C1 flare. As GOES and XSM are essentially measuring the same thing, and the measurements have been calibrated for the detector responses, then ideally the ratio should be equal to 1 (with an -21-