Chandra Observations of Comet 9P/Tempel 1 During the Deep
Impact Campaign
C.M. Lisse1, K. Dennerl2, D. J. Christian3, S. J. Wolk4, D. Bodewits5, T.H. Zurbuchen6,
K.C. Hansen6, R. Hoekstra5, M. Combi6, C. D. Fry7, M. Dryer7,8 , T. Mäkinen9, W. Sun10
Submitted to the Journal Icarus, 17-Apr-2006; accepted 09-Mar-2007
1
Planetary Exploration Group, Space Department, Johns Hopkins University Applied Physics Laboratory, 11100
Johns Hopkins Rd, Laurel, MD 20723, carey.lisse@jhuapl.edu
2
3
Max-Planck-Institut für extraterrestrische Physik, Postfach 1603, 85740 Garching, Germany, kod@mpe.mpg.de
Queens University Belfast, Department of Astronomy and Physics, Belfast, UK BT7 1NN D.Christian@qub.ac.uk
4
Chandra x-ray Center, Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA
swolk@cfa.harvard.edu
5
KVI atomic physics, Rijkuniversiteit Groningen, Zernikelaan 25, 9747 AA Groningen, the Netherlands,
bodewits@kvi.nl
6
The University of Michigan, Department of Atmospheric, Oceanic and Space Sciences, Space Research Building,
Ann Arbor, MI 48109-2143, USA
7
Exploration Physics International, Inc., 6275 University Drive NW, Suite 37-105, Huntsville, AL 35806-1776,
gfry@expi.com
8
NOAA Space Environment Center, 325 Broadway, Boulder, CO 80305, Murray.Dryer@noaa.gov
9
Finnish Meteorological Institute, Space Research, P.O. Box 503, FIN-00101 Helsinki, Finland,
teemu.Mäkinen@fmi.fi
10
Geophysical Institute, University of Alaska, Fairbanks,AK 99775, sun@jupiter.gi.alaska.edu
36 Pages, 7 Figures, 4 Tables
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Proposed Running Title: Chandra-DI Observations of Comet 9P/Tempel 1
Please address all future correspondence, reviews, proofs, etc. to :
Dr. Carey M. Lisse
Planetary Exploration Group, Space Department
Johns Hopkins University, Applied Physics Laboratory
11100 Johns Hopkins Rd
Laurel, MD 20723
240-228-0535 (office) / 240-228-8939 (fax)
Carey.Lisse@jhuapl.edu
2
Abstract
We present results from the extensive Chandra X-ray Observatory’s campaign of Comet
9P/Tempel 1 (T1) in support of NASA’s Deep Impact (DI) mission. T1 was observed for ~300
ksec between 30th June and 24th July 2005, and continuously for ~60 ksec on July 4th during the
impact event. x-ray emission qualitatively similar to that observed for the collisionally thin
comet 2P/Encke system (Lisse et al. 2005b) was found, with emission morphology centered on
the nucleus and emission lines due to C, N, O, and Ne solar wind minor ions. The comet was
relatively faint on July 4th, and the total increase in x-ray flux due to the Deep Impact event was
small, ~20% of the immediate pre-impact value, consistent with estimates that the total coma
neutral gas release due to the impact was 5 x 106 kg (~10 hrs of normal emission). No obvious
prompt x-ray flash was seen. Extension of the emission in the direction of outflow of the ejecta
was observed, suggesting the presence of continued outgassing of this material. Two flares,
much stronger than the man-made increase due to Deep Impact, were found in the observed xrays on June 30th and July 8th, 2005, and are coincident with increases in the solar wind flux
arriving at the comet. Modeling of the Chandra data using observed gas production rates and
ACE solar wind ion fluxes with a CXE mechanism for the emission is consistent with the
temporal- and spectral behavior expected for a slow, hot wind typical of low latitude emission
from the solar corona interacting with the comet's neutral coma.
Key words: Comets, Comets Individual: 9P/2005 (Tempel 1), Atmospheres, Coma, Solar Wind,
Spectroscopy, X-rays
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I
Introduction
Since the surprise discovery of x-rays from comets in the mid-1990s using ROSAT and EUVE
(Lisse et al. 1996, Mumma et al. 1997, Dennerl et al. 1997, Krasnopolsky et al. 2000) over 20
comets have been detected in the x-ray regime, including many with the most advanced x-ray
observatories, Chandra, CHIPS, and XMM-Newton (Lisse et al. 2001, Krasnopolsky et al. 2002,
Dennerl et al. 2003; Wegmann et al. 2004; Lisse et al. 2005a,b; Sasseen et al. 2006; Bodewits et
al. 2007). These results have further refined our understanding of cometary x-ray emission via
charge exchange between instreaming solar wind minor ions and outflowing gaseous neutral
cometary material. However, these all have been passive observations. NASA’s Deep Impact
(DI) mission changed all that, by impacting comet 9P/Tempel 1 (hereafter T1) with a copper
impactor on 04 July at 05:44 UT at a relative velocity of 10.2 km sec-1. The collision of the 364
kg impactor released 106 – 107 kg of fresh neutral cometary material into the extended
atmosphere above the comet (A’Hearn et al. 2005; Meech et al. 2005, Schultz et al. 2007),
increasing the neutral gas mass of the coma temporarily by up to 25%, and providing a
controlled, highly monitored experiment to observe possible changes to the x-ray emission.
Knowing exactly when this new neutral material arrived allows us to test the charge exchange
emission mechanism (CXE; Cravens 1997, Lisse et al. 2001, Kharchenko and Dalgarno 2001,
Kharchenko et al. 2003, Beiersdorfer et al. 2003, Bodewits et al. 2006). Because of the comet’s
low outgassing rate (Schleicher et al. 2006), comparable to that of 2P/Encke (Lisse et al. 2005b),
it was collisionally thin to charge exchange. The new material should therefore have directly
increased the emission measure for the comet, and possibly changed the signature of the CXE
emission to a collisionally thick system. Changes in the solar wind flux at the comet also cause
changes in the observed x-ray flux, and need to be carefully accounted for. One of the other
exciting aspects of the impact was to see if enough energy was released to create prompt x-rays
by hyper-velocity impact processes, as suggested by ROSAT observations of the impact of the
K-fragment of comet D/Shoemaker-Levy 9 on Jupiter (Waite et al. 1995).
In this paper, we report the results of our extensive Chandra x-ray Observatory (Chandra)
observing campaign of T1 during the Deep Impact encounter. The Chandra observations, in
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addition to showing possible changes in the x-ray emission as the result of the impact, provide
one of the longest duration studies of changes to cometary emission in the x-ray. We describe the
details of the Chandra and ACE instruments used, observing strategy, and our data analysis
techniques in §II. In §III we discuss the solar wind environment during the Chandra
observations and the comet's outgassing behavior. In §IV we present the Chandra observational
results, including T1’s morphology, temporal changes and spectra, and place these in the context
of previous cometary x-ray observations. We compare changes in the x-ray emission to the
behavior of the solar wind and coma, and discuss the most dominant processes responsible for
the observed variations. We then discuss the changes caused by the Deep Impact experiment. In
§V we summarize our findings, the most plausible physical model for Tempel 1, and what was
learned from the x-ray observations of the DI experiment.
II
Instruments and Observations
A. Chandra x-ray Observatory. The Chandra observations and data analysis were done in the
same manner as for our previous comet observations (Lisse et al. 2001, 2005b) and the reader is
referred to these papers for details of the observational scheme. Briefly, the ACIS-S CCD was
used, providing high resolution x-ray imaging with platescale of 0.5"/pixel, instantaneous field of
view 8.3' x 8.3', and moderate resolution spectra (E ~ 110 eV; FWHM, Gaussian = 50 eV) in the
300 to 2500 eV energy range. The ACIS-S optical axis was pointed on the sky 2' in front of the
comet along the direction of its apparent motion, and the exposures were taken as the comet
moved through the ~4' field of view. T1 was centered in the back-illuminated CCD chip S3, the
most sensitive CCD for energies below 1 keV. No active guiding on the comet was attempted,
and Chandra was able to follow the comet's non-sidereal motion using multiple pointings.
Chandra observations of comet T1 were planned to maximize our detection of any changes
caused by the impact. For this reason preliminary observations were obtained on June 30th to
establish the baseline emission of the comet, and long duration observations were obtained
immediately before, during, and after the impact on July 4th and 5th. These observations were
followed by several monitoring “snapshot” observations on July 8, 10, 13 and 24. (Table 1).
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Each pointing obtained between 28 and 64 ksec on the comet for a total integration time of ~295
ksec.
Combined images from each set of observations were re-mapped into a coordinate system
moving with the comet using the “sso_freeze” algorithm, part of the Chandra Interactive
Analysis of Observations (CIAO) software. Source photons from the comet were extracted from
a circular region of 3.7' centered on the comet, with a background region defined from the outer
part of the CCD between 4' from the nucleus and the edges of the S3 chip. Images, light curves
and spectra were extracted with the CIAO tools and analyzed with a combination of IDL and
FTOOL packages. The x-ray spectral analysis software package XSPEC (Arnaud 1996) was
used for the spectral fitting.
B. Advanced Composition Explorer
The Advanced Composition Explorer (ACE) carries solar wind plasma and composition
instruments that, since its launch in late 1997 have provided the best coordinated solar wind
dynamics and composition measurements, including elemental, ionic and isotopic abundances
and fluxes of the relevant particle species found in interplanetary space (Stone et al. 1998). ACE
is located in a halo orbit near the Lagrangian libration point L1. Due to its real-time capability
for space environment data, ACE data are being used for real-time space weather forecasting
world-wide (Zwickl et al. 1998).
The instrument of main interest for this study is the Solar Wind Ion Composition Spectrometer
(SWICS), an ion mass spectrometer that resolves the elemental and ionic composition of all
significant solar wind species. SWICS is a combination of three key systems: an electrostatic
analyzer, a time-of-flight telescope, and a solid state detector array (Gloeckler et al. 1998).
SWICS resolves solar wind ion species by measuring, at the same time, solar wind speed, mass
per charge ratio with ~5% resolution, and ion mass with a resolution of the order 15-20%. The
larger the mass of the solar wind ion detected, the coarser the mass resolution. The analysis of
SWICS data therefore uses a process separating neighboring and overlapping peaks (von Steiger
et al. 2000). This process does a very good job analyzing the most abundant ion species and
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charge states, but has limits when analyzing very small fluxes of low-abundance charge states
near multiple large peaks, e.g. the extremely rare but CXE important ion N7+.
III. Solar Wind and Comet Characteristics at the Time of the Chandra
Observations. The flux and the heavy ion composition of the solar wind at the comet, and the
comet's outgassing rate both strongly affect the observed x-ray emission (Lisse et al. 1999,
2001). A detailed analysis of the solar wind dynamics, composition/charge state, and the comet's
activity are therefore a crucial part of our analysis, and we go into some detail describing them
here.
IIIA. Large Scale Modeling of the Solar Weather.
The nearly 30 day time span of the Chandra observations contained several significant changes
in the solar wind heavy ion and proton environment (Figures 1 and 2, respectively). To
understand the large scale solar wind interplanetary structures influencing the comet's x-ray
emission behavior, we utilize here the well tested Hakamada-Akasofu-Fry solar wind model
(HAFv2; Figure 2, Table 2), which takes into account the source solar coronal structure
behavior, and follows the solar wind propagation through interplanetary space while allowing for
solar wind plasma density, temperature and velocity dispersion relationships. We only quote the
results of using the HAFv2 model here; we refer the reader to Fry et al., 2001, 2003, 2004;
Dryer, 1998; and Dryer et al., 2001, 2004 for more details of the model.
On June 30, at the beginning of the period, the solar wind was dominated by a slow stream at
~350 km/s. This condition quickly changed on the latter part of June 30th as the heavy ions’ wind
speeds rose to 650 km/s (Figure 1) while maintaining roughly the same density. Within 3 days
after the Co-rotating Interaction Region’s (CIR’s) passage, as suggested by the compressed
Interplanetary Magnetic Field (IMF) lines in Figure 2, the solar wind ramped down in a smooth
fashion to the pre-June 30th state. On July 4, Tempel 1 was about 35º east of the Earth-Sun line
essentially in the ecliptic. The western flank of the modeled shock in Figure 2 (from a July 1 st
C5.3 solar flare at N14E83 as noted as FF#608 in Table 2) might have intercepted the comet on
July 4th with a somewhat higher flux of proton and heavier ions (C, N, O, Ne). We expect,
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however, any solar wind signature on July 4th from the above-mentioned shock to have been
relatively weak (Figure 2).
There was a significant shock (FF#611, Table 2; see caption for Figure 2) from another solar
flare on July 9th at N19W28 accompanied by a fast (~1600 km/s at the Sun) halo Coronal Mass
Ejection (CME). The shock appears in both our model and in the ACE solar wind proton
measurements (Figure 1, Sec III.B). The further, weakened eastern flank (see Figure 2’s caption)
of this flare shock passed over Tempel 1 on July 13th (DOY 194), causing milder changes in the
observed x-ray behavior. What makes this event very different than the June 30 th event is that
the accompanying Interplanetary Coronal Mass Ejection (ICME) produced prompt, hot ions
moving rapidly and radially from the Sun, widening and prolonging the solar wind perturbation
before the much slower shock reached the comet. (A similar situation with a radially traveling
ICME was found for comet C/1999 S4 (LINEAR) in 2000 (Lisse et al., 2001)). This led to a
prolonged increase in the solar wind flux at the comet and ACE, starting on July 9 th and
producing a prolonged x-ray response. Finally, it should be noted that the solar flares in Table 2,
from FF #612 onward, were likely ineffective in “reaching” T1 because of their far-westward
and solar backside locations. Instead, CIRs were the major global solar wind structures passing
over T1 in our case, thereby justifying the mapping procedure discussed earlier. We note that
continuous modeling is essential to gain a global temporal and spatial perspective of the highly
variable heliosphere with both co-rotating and transient structures, the latter considerably altering
the former.
IIIB. ACE Measurements During the Chandra Observations. ACE, the source of all our
major solar wind data, was located at roughly 34-42o west of the comet (as observed from the
Sun) during the observations. Analysis of the ACE solar wind data has to be done in two
successive steps: In the first step, the calibrated heavy ion fluxes are computed from the
spacecraft telemetry and the resulting temporal trends are produced (Figure 1a). The second step
in the determination of the solar wind ion flux at the comet is the mapping of the ACE data from
the near-Earth ACE orbit to the comet’s heliocentric location. During the Chandra observations,
the comet was essentially in the ecliptic plane and therefore had a negligible latitudinal
separation from ACE (see Table 1). This allowed us to use a 1-D MHD mapping as described in
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Prange et al. (2004). This process takes into account an expansion of the solar wind, and its
dynamic evolution during this expansion. Due to the relatively modest, ~0.5 AU expansion, all
dynamic signatures seen in Figures 1 and 2 remain visible. We then assume persistence in the
solar wind structure during 35o of solar rotation. (This is equivalent to 2.5 days of a rotation of a
Carrington spiral, so our assumption is valid for the large-scale solar wind expansion, but may
not be very good for CMEs.) The middle panel of Figure 1 shows the solar wind speed patterns
after the detailed radial calculation, shifted to account for the longitudinal offset. The simulated
solar wind conditions are discussed in Fig. 2.
Overall, we see a tremendous degree of variability of heavy ion densities in the ACE data, due
mainly to three different effects. (1) There are dynamic changes in the solar wind. E.g., there are
a number of fast velocity changes caused by transitions from a slow/coronal streamer associated
wind, to a fast/coronal hole associated wind. Such transitions occur near July 2, 25, and 31 (DOY
183, 206, and 212) and are labeled A, B, and C in the figure, respectively. All speed increases
also lead to large increases in temperature, and hence the thermal speed of heavy ions. Some of
these transitions are shock-like, but other compressions have not succeeded to form shocks at the
time of the observation. There is another abrupt increase in speed on July 10 (DOY 191, denoted
D) and a less dramatic increase on July 18 (day 199, denoted E). Both of them are different in
nature and associated with either a co-rotating interaction region (CIR) or an Interplanetary
Coronal Mass Ejection (ICME, a high velocity, transient solar flare-generated CME preceded by
its low corona-formed shock wave). There is a long-term persistence of fast solar wind streams
(formed by high speed flows from coronal holes that compress preceding slower flow, thereby
forming shocks), lasting many days, as demonstrated by Smith and Dryer (1991). (2) The heavy
ion compositions can change due to variations in elemental abundance. The variability of the
elemental abundances is of the order of a factor of 2 (von Steiger et al. 2000, and references
therein), and are therefore a minor part of the variability shown here, which is over 2 orders of
magnitude. We see such effects on July 10 (DOY 191). (3) The solar wind (SW) heavy ion
composition changes due to the electron temperature of the solar wind sources (see Zurbuchen &
Richardson 2006). These variations have shown to be the best signatures of CMEs and coronal
holes, and are clearly visible in the top of Figure 1. There are correlated decreases of O 7+ and
N6+, associated with coronal hole winds, on July 2, 25, and 31 (DOY 183, 206, & 212). CME
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dominated winds, as seen on June 30 (DOY 181) do not show such decreases, but instead lead to
significant increases of high charge states, such as O8+ (and presumably also N7+).
In summary, we expect from the HAFv2 models, and see from the ACE data, large changes in
the solar wind fluxes and charge states at the comet on June 30, July 9 - 10, 13, 18, 25, and 31 of
2005.
III.C Measures of Tempel 1's Outgassing Behavior. Unlike our previous Chandra comet
observations, the comet was being heavily monitored at other wavelengths sensitive to the
neutral gas emission rate from the nucleus. The solar wind and cometary dust and gas emission
activity were all studied in depth and length, by over 80 observatories around and above the
Earth. In fact, one of the key science results to come from the Deep Impact experiment is the
verification of the CXE mechanism as driving cometary x-ray emission using multiple lines of
evidence.
The Deep Impact spacecraft observed the comet's behavior in detail from 30 June 2005 - 5 July
2005. A number of brief, but large explosive outbursts were seen by the DI and also by the HST
spacecraft over the period of June - July 2005 (A'Hearn et al. 2005, 2007; Farnham et al. 2007).
Rise times for these outbursts were typically on the order of minutes, and fall times typically on
the order of hours. A direct comparison of the DI detected outbursts and the Chandra detected
outbursts does show a correspondence for the "weak" outburst of 30 June 2005 with the strong
Chandra outburst of the same day. There is no DI data corresponding to the 8 July 2005 strong
Chandra outburst, and observations by other observatories do not show enhanced neutral gas
emission from the comet on this day.
Lowell Observatory observed the production rate of OH, the direct daughter product of H2O, the
majority neutral gas species emitted by the comet, over the time span June 29 to July 11. The
Spitzer Space Telescope monitored the total 16 m brightness, due to thermal continuum
emission from dust emitted by the comet, over the time span July 2 - July 14, 2005. Correcting
both of the water daughter product datasets for the mean lifetime of water vs. UV photolysis at
1.51 AU, ~ 5 x 104 seconds, we find good agreement with the relative Spitzer dust production
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temporal trends where they overlap in time. Here we use the combined curve as a scaled proxy
for the production rate of neutral gas molecules at the nucleus, Qgas.
Corresponding measurements of the Ly- emission from neutral H-atoms over the time range 5
to 30 July by the SWAN UV instrument onboard the SOHO spacecraft (Mäkinen et al. 2001,
2007) was obtained as a check of our method. The neutral H-atom coma of a comet is produced
by the same UV photolysis of water that produces OH. While sparsely sampled, the
SOHO/SWAN lightcurve is taken by a stable spacecraft platform, is long in duration, and has
been calibrated and verified for numerous comets. The water production rate found by SWAN
(see inset for Fig 4a) was flat from 5 July through 10 July, in agreement with our derived Qgas
curve. An apparent upturn, of amplitude 50 - 150 % and uncertain duration, was seen on July 14
- 15, with the production rate returning to the average quiescent value of ~1 x 1028 mol sec-1 by
July 17.
IV.
Chandra Results
IVA. Morphology. Chandra ACIS-S3 images of the 9P/Tempel 1 x-ray emission are shown
in Figure 3. T1’s emission morphology is centered on the nucleus and is nearly point-like, unlike
the majority of the 20+ comets observed in the x-ray to date (Lisse et al. 2005a). The point-like
morphology is very similar to the morphology seen for 2P/Encke 2003 (Lisse et al. 2005b),
which was determined to be due to a collisionally thin coma, a result of the low rate of neutral
gas emission from the nucleus. The reported rate of water gas production for T1, 1 to 3 x 10 28
mol s-1, was very close to the 2 x 1028 mol s-1 for 2P/Encke in 2003. In the collisionally thin
regime, the maximum of the comet's emission occurs very near the nucleus, where the maximum
density of coma neutral gas is found. With T1 at 0.88-0.98 AU from the Earth, 1" corresponded
to 630-710 km at the nucleus, so that the ACIS-S3 pixel (0.5" x 0.5") scale was 315-355
km/pixel at the comet.
However, the actual size of a statistically significant (i.e. 5) spatial resolution element was
determined by the limited number of photons detected. In regions of high net signal, e.g. near the
nucleus, this resolution was on the order of 2" (~1300 km) per resolution element. The average
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photon limited effective pixel size realized over the detected x-ray emission region was 7"
(~4700 km) per resolution element. We find a maximum extent of detectable x-ray emission of
~200,000 km for the comet. This implies an x-ray extent larger than the bowshock distance:
according to the formula of Wegmann (2004), rbowshock = 1.7 x 104 km June 30; 4.5 x 103 km on
July 4 before impact and the other quiescent periods; and 1.0 x 104 km on July 8. As the comet
was weakly active (Qgas ~ 1028 mol sec-1, Makinen et al. 2007), we can expect penetration of the
SW to near the nucleus. Using the results Giotto Grig-Skellerjup flyby (Johnstone et al. 1993,
Lebedev 2000), we estimate a radius for the SW ion contact surface, where molecular collisions
become dominant in the coma, of < 100 km, much less than one Chandra ACIS-S3 pixel in
extent.
The strongest evidence for deviations from spherically symmetric gas emission behavior from
the nucleus occurred in our observations on the days of 30 Jun 2005 and 4-5 July 2005 (DOY
181 and 185-186). On June 30th, a brightening was seen in the anti-solar direction, increasing the
radial extent of the emission by a factor of 50 ± 25%. We interpret this as evidence for an
outburst of neutral material from the comet with projected direction in the anti-solar direction.
On 4-5 July 2005, an extension was seen in the images in the same SW direction as the flow
observed in near-nucleus images of the ejecta produced by the impact (A'Hearn et al. 2005,
Meech et al. 2005, Sugita et al. 2005). The statistical significance of the extension is ~5 above
background. Larger scale imaging by the ROSETTA spacecraft over the course of the next few
days confirms the progression of the water gas and OH ejected by the impact (Kueppers et al.
2005). The estimated propagation velocity of the extension is 0.5 to 1 km sec-1 for the ~33hr
duration of our 4-5 July observations, but we are limited by total signal and possibly viewing
geometry. The derived expansion velocity is higher than the ~0.2 km sec-1 found for the leading
edge of the solid ejecta (Lisse et al. 2006, Schleicher et al. 2006), and we associate it instead with
the higher expansion velocity of neutral gaseous material produced by the impact, as expected
for CXE driven x-ray emission. We use the estimated total mass of water ejected by DI (~5x106
kg), and our estimate of a 10x-ray detection by Chandra over 4-5 July in a total on-target time
of 130 ksec to estimate the limiting sensitivity of Chandra to cometary activity as 5x106 kg/130
ksec/10 = 4 kg sec-1 (1, equivalent to ~1026 mol sec-1 of water). Comparison of the Chandra x12
ray morphology to the ROSETTA OH images shows that x-rays are not the most sensitive way
to detect OH neutral gas; optical narrow band imaging has a superior emitted flux/background
statistic.
IVB. Light Curve. A Chandra x-ray light curve was created for Tempel 1 using a 300 –
1000 eV passband, with corrections applied for the 24% increase in projected area from June 30
to July 24. The 300 – 1000 eV range allows the best tradeoff between exclusion of instrumental
background and inclusion of maximal comet signal (Figure 4a). Impulsive events were found in
our data and these are similar to previous observations of comets Hyakutake, Encke, and Linear
S4 (Neugebauer et al. 2000, Lisse et al. 1999, 2001). Overall, for an average distance to T1 of
1.507 AU, the 300 – 1000 eV x-ray average luminosity, Lx was 4.6 x 1014 erg s-1. As comet
Tempel 1 was collisionally thin to charge exchange to within a fraction of the central pixel, the
x-ray lightcurve is directly proportional to the comet’s neutral gas emission activity, the flux of
solar wind minor ions in the coma, and the cross section for CXE between the two. The cometary
x-ray emission observed is driven by the relatively abundant fully stripped, hydrogen-like, and
helium-like C,N,O, and Ne ions. All Lyman- or K-series photons emitted after electron capture
by these ions are within the energy range detectable by Chandra (0.3-1.0 keV), and therefore the
total cross sections for photon emission are equal to the total charge exchange cross sections
(Bodewits et al. 2007). As high signal-to-noise measures of C,N,O, and Ne ion flux
measurements are not always available for use in our analysis, we use both the direct C,N,O and
Ne measurements where they are available, and the more frequent and higher signal-to-noise
proton flux measurements to bridge the gaps in coverage, assuming they trend in the same
overall fashion as the minor ion fluxes.
The long term Chandra lightcurve is shown in Figures 4a, along with the time history of the
comet’s gas production rate estimated from Lowell narrowband OH filter measurements
(Schleicher et al. 2006) and Spitzer continuum data (Lisse et al. 2006) and the time history of
ACE/SWICS solar wind proton flux data. Qgas has been estimated using the QOH outgassing rates
from Schleicher et al. (2006), shifted by -1 day, to allow for the ~105 seconds required to
photolytically form OH from the parent water molecule emitted by the nucleus, combined
13
directly with the 16 µm Spitzer lightcurve data. (The Spitzer data are dominated by thermal
emission from coma dust, and thus measure changes in coma emission activity promptly,
assuming the dust and gas emission rates from the nucleus are coupled.) This assumption seems
valid, as the Lowell and Spitzer trends agree very well over the 8 day overlap period 2 to 10 July
2006. The predicted lightcurve formed from the product of Qgas and the wind factor is also
shown. The observed x-ray fluxes mirror the solar wind behavior well, except on June 30th (DOY
181), and at impact on July 4th (DOY 185). Tempel 1 also represents the first comet to be
observed and detected simultaneously by multiple orbiting x-ray telescopes (Chandra, SWIFT,
and XMM). In Figure 4a we have included the SWIFT photometric data for the comet
(Willingale et al. 2006), which agrees well with the trends observed by Chandra.
On June 30th we find that the measured solar wind flux increase can account for at most 50% of
the observed x-ray flux increase; the rest must come from increases in the comet's neutral gas
outflow rate. A corresponding increase in brightness was indeed found in the Deep Impact
approach lightcurves for this date (Farnham et al. 2007). This was followed by a period of
relatively quiet wind and low outgassing around the time of impact.
On July 4th we examined the observed x-rays for any direct indication of changes as the result of
the impact. While there is a possible prompt spike jump in the very first points of the lightcurve
measurements after impact, it is only found after smoothing the data with effective bin sizes of
5000 s. We thus do not claim a detection of an impact flash in the x-ray. We estimate the
maximum amplitude of any flash to be less than 0.02 cps (30% of the pre-impact quiescent count
rate). Using this upper limit, we estimate that it would take at least 5 times as much signal, and
hence at least 5 times as much energy deposited onto a cometary body at a distance of 0.9 AU
from Chandra, to make a significant x-ray flare we could detect.
We also searched the July 4 light curves for any possible enhancements immediately after impact
and also searched for any decreases in emission in the interval following the impact, assuming
either absorption or quenching were possible. We find we can rule out changes in the solar wind
as the cause of the small, but significant (18 ± 5%) increase found in the average trend of the xray emission over the course of July 4. This emission rate stayed high on July 5, although the
14
solar wind flux trends did not show this level of increase. The increase in the observed x-ray flux
on July 4th is consistent with a total amount of water gas released by Deep Impact of ~5 x 10 6
kg, as reported from ROSETTA measurements of coma OH emission during and after the impact
(Keller et al. 2005, Küppers et al. 2005). This amount of water gas, equivalent to approximately
10 hours of the normal ambient outflow rate of about 1 x 1028 mol s-1, is roughly 30% of the total
gas contained in a coma of radius 1 x 105 km.
The July 8th (DOY 189) flare appears to be driven mainly by a strong, prolonged increase in the
solar wind flux (Section III). After another period of relative quiescence, the comet again
encountered increased solar wind fluxes on July 16 (DOY 197), which had disappeared by July
21. By the time of the very last Chandra observation on July 24, the solar wind flux had begun to
increase once again.
It is important to note here that there is no evidence for the large upsurge in material (1.9±0.4 x
108 kg of water or 3.9 ± 0.5 x108 kg of carbon dioxide) emitted from the comet into the coma due
to the DI experiment as reported by SWIFT, starting with July 8 (Willingale et al. 2006); instead,
the broad increase in x-ray flux seen in the combined Chandra/SWIFT x-ray lightcurves is
explained by increases in the solar wind flux at the comet (Figure 4a). There is also no obvious
increase in the SWIFT x-ray count rate during the increased outgassing on June 30th - July 1,
2005. The long term temporal behavior of the x-ray luminosity, before, during, and after the
Deep Impact experiment, can be entirely understood using measures of the coma neutral gas
density (Küppers et al. 2005, Jehin et al. 2006, Lisse et al. 2006, Schleicher et al. 2006) that
show an elevated density of material for the first two days after impact, then normal cometary
behavior afterwards.
Nucleus rotation effects. For 2P/Encke, we found a lightcurve responding to diurnal changes in
outgassing due to the rotation of the asymmetric nucleus with discrete regions of enhanced
outgassing (Lisse et al. 2005). Because the comet was collisionally thin to CXE, we expected
and found prompt response to increases in coma neutral density, although our detection was
somewhat noisy. Here we have attempted a similar experiment to confirm this result : taking the
ensemble of Chandra data and searching for the best-fit long term lightcurve, to see if it is
15
consistent with other measures of activity modulation due to nucleus rotation (e.g., the Deep
Impact spacecraft approach imagery lightcurve (A'Hearn et al. 2005) or the ground based CN
narrow band photometry of Jehin et al. (2006)). Allowing for effects due to the windowing
function, episodes of cometary outburst, episodes of increased solar wind flux and artifacts due
to the fixed pointing + comet drift method, we find that phasing the data to a period of 1.70 days
(as reported by Lisse et al. 2005c, Jehin et al. 2006, Belton et al. 2007, Farnham et al. 2007)
yields a periodic waveform. Running a 2 period search, we find the best fit period to be 1.6 ±
0.15 days (Figure 4b). Jehin et al. show a T1 folded CN/NH lightcurve from 3-7 July 2005 with
peak-to-peak amplitude of 50% to have the highest peaks between phases 0.0 and 0.3, similar to
the maxima at phases of 0.1 - 0.4 in the Chandra data.
Removing the high count rate flare data of June 30th and July 8th from the fit, however, reduces
the significance of the periodic fit substantially. As the flare data represent only 20% of all
photons, this change should not be large, but a simple calculations shows that the flares are
separated by exactly 5 rotations if we assume a 1.63 day period. We thus consider the rotational
detection in the x-rays to be somewhat weak, and report here that the x-ray lightcurve is
consistent with a 1.6 - 1.7 day period.
IVC. X-ray Spectroscopy. The Chandra 9P/Tempel 1 spectra demonstrated the classical
emission features of cometary x-ray emission (Figure 5), with strong emission due to highly
charged C, N, O, and Ne solar wind ions charge exchanging with neutral gas species in the
cometary coma (Lisse et al. 2001, 2005a; Krasnopolsky et al. 2002). In modeling the faint T1
comet spectrum, we have required our spectral model to be consistent with current models of
charge exchange emission (Bodewits et al. 2004) and models of the observed spectra of 7 other
comets measured by Chandra (Bodewits et al. 2007). We have done this by fixing the line
energies of the model to the C, N, O, and Ne emission lines detected in the other comets,
effectively minimizing the number of free parameters in our line emission models. While we
found that other spectral models, such as the multiple emission line + thermal Bremsstrahlung
continuum (Lisse et al. 2001) could also produce acceptable fits to the data, we argue that the
observed “continuum” could be understood to be the blending of many weaker lines, and is not a
self-consistent or physically viable model.
16
Specifically, we have combined the theoretical predictions of the CXE model of Bodewits et al.
(2004, 2006) with a 1-D Haser model of the neutral coma gas density (Qneutral = QH2O + QOH +
QCO2 + QCO + QH + QC + QO; Figure 6). The Bodewits model provided the line energies in the
300 to 1000 eV energy range from solar wind minor ion species CV, CVI, NVI, NVII, OVII,
OVIII and Ne IX. Our CXE model has 34 lines, but we included only the 11 most significant in
our modeling (and no continuum): those based on the strongest lines expected from the solar
wind composition for our CXE model, and also the most statistically significant lines from fitting
all spectra also be those first determined from the very strong June 30th flare spectrum. The data
were then fit at these energies in XSPEC as a series of Gaussian emission lines, each at a fixed
line energy and set to the ACIS-S3 instrument resolution. Model fits from all 7 good T1 spectral
observations were compared to determine the significant emission lines in a consistent manner.
The strongest emission features in the T1 spectrum are due to the OVII multiplex between 550 580 eV. Other relatively strong emission features found were the NVII/CVI and CV lines
between 350 and 460 eV. Faint lines from NVII (500 eV) and Ne IX (922 eV) were also
detected. This is the first time NVII lines have been definitively detected, and only the 3rd comet
for which the NeIX lines were found. Spectra observed during epochs of intensified solar wind,
namely June 30, July 8th and July 24th also showed an additional line at ~605 eV. This line may
possibly be NVII Ly- transition at 595 eV, but from the CXE model this line is expected to be
weak and only ~10% of the NVII 500 eV Ly- emission. We present results of the best fitting
multi-emission line model in Table 3. The number of degrees of freedom of the fit, equivalent to
the number of independent spectral channels in the Chandra observations minus the fit
parameters, was 36. NH, the neutral hydrogen column density between Chandra and the comet,
was set to 1013 cm-2 (i.e. effectively to zero compared to typical columns of 1019 cm-2 or greater
found in astrophysical x-ray sources) for the models, as expected for the density of hydrogen in
interplanetary space.
The June 30th flare spectrum has over 150 counts per channel in the OVII line and this
spectrum was of good signal-to-noise. Arguably, several of the fainter spectra (July 5, 10th)
can be modeled using only 6 or 7 emission lines, but we have included all lines from our CXE
17
model to compare spectra of the different observations and for comparison to previous results
in the literature (e.g Lisse et al. 2001, 2005; Krasnopolsky et al. 2002; Kharchenko et al. 2003)
The large errors (or upper limits) of these fainter lines are reflected in the summary of fitting
results in Table 3,
We note that an additional emission line with energy 270 to 285 eV could be included in the
spectral models. However, the Chandra effective area in this region is low because of the
instrument’s carbon edge. This fact, coupled with the rising x-ray background makes the
emission features below 300 eV suspect and we limited the spectra to start at above this region.
This limitation did not cause us to lose any line identifications, while making our modeling more
robust by increasing the measured S/N ratio.
X-ray Spectra and Charge State Ratios. Our Chandra observations sampled several very
different solar wind environments (Sec. IIIA,B) and thus provide at least two different spectral
signatures with which to test the CXE model. Because the comet remained collisionally thin
throughout the observations, compositional variations in the solar wind should directly result in
spectral variations. Charge exchange emission from H-like ions is dominated by strong Ly-
emission, or, equivalently for He-like ions, by the forbidden-intercombination-resonance feature
(FIR). From state selective electron capture cross sections we derived velocity dependent
emission cross sections for all emission lines of interest. These cross sections allow us to extract
relative ion abundances from the observed spectra based on the proton speeds derived from the
ACE data (Bodewits et al. 2007). Comparison of the ion abundances derived from the x-ray
emission to those measured by ACE and propagated back to the comet are given in Table 4 and
Figure 5b. As discussed above, the solar wind was highly variable during the T1 observations,
and this resulted in large variations in the observed abundances for C, N , O. For carbon and
oxygen, the obtained values agree within the uncertainties with averaged slow solar wind ion
abundances, such as those listed in Schwadron & Cravens 2000 (0.35 and 1.59 for O8+ and C6+
relative to O7+, respectively). The C6+ ion ratio appears to track the ACE SW data especially
well.
18
The carbon ion, C6+, has the best statistical accuracy of the three ions presented in Figure 5, and
demonstrates the best agreement with the Schwadron & Cravens (2000), Beiersdorfer et al.
(2003) and Otranto et al. (2006) SW trends (see Table 4). The derived N7+ abundance (0.03
relative to O7+), however, is consistently high by almost an order of magnitude as compared to
average solar wind values (Schwadron & Cravens 2000). The O+8 abundances are less
problematic, but still high by about a factor of 2 compared to the Schwadron & Cravens 2000
estimates and the analysis results of Chandra observations of C/McNaught-Hartley by
Krasnopolsky et al. (2002).
While at first apparently alarming, careful study of our methodology has not been able to find an
error in our spectral modeling that could account for these discrepancies - the required errors in
the CXE cross sections are too large compared to the uncertainty of the measurements
(Beiersdorfer et al. (2003) and Otranto et al. (2006)). Other issues may be in play. E.g., our work
with solar wind trending data (von Steiger et al. 2000, Zurbuchen and Richardson 2007) has
determined that the abundance ratios of the relatively rare ions used in Schwadron & Cravens
2000, such as O8+ and N7+, have not been determined to better than an order of magnitude to date
(due to poor counting statistics). Also, the Schwadron & Cravens 2000 values for average solar
wind were derived on a time-interval in a very different part of the solar cycle, with different
average compositions. Further, the minor ion abundances in the SW are highly variable, so that
the value of <average SWabundance>/RMS SWabundanc rarely exceeds the value of 1 over the course
of a day (see Figure 1). Given the faintness of Tempel 1 in the x-ray, however, and the relatively
large uncertainties of our ion charge state estimates (Figure 5b, Table 3) we do not feel the issue
can be solved with the present dataset. We raise the issue here as one requiring further study,
which is now possible, using spectral analyses of other, brighter comets observed by Chandra
(Krasnopolsky 2006, Bodewits et al. 2007).
Overall, the charge state of the solar wind encountered by the comet was close to the relatively
hot, dense, slow wind emitted from the Sun's lower latitudes around the minimum in the solar
activity cycle. This stands in strong contrast to Chandra results of comet 2P/Encke on 24 Nov
2003 (Lisse et al. 2005b), where we measured a very unusual x-ray spectrum with approximately
equal strength CVI and OVII CXE lines in the 300 - 700 eV range during a period of mixed slow
19
and fast winds. In fact, the new Tempel 1 results, for another collisionally thin system, now
allows us to empirically separate behavior induced by strong changes in the solar wind incident
charge state from that due to a collisionally thin coma. Tempel 1 showed a "normal" spectrum
(i.e., a ~ 3:1 line ratio of the 560 eV/ 400 eV complexes) for most of the time it was observed.
reassuringly, however, the spectrum of July 5th, 2005 appears to be rather like that seen for
Encke in 2003.
Pre- and Post- Impact Spectral Behavior of Tempel 1. Because of the flare during
the June 30th observation, which had been intended by our experimental design to establish the
quiescent baseline value of the comets emission, our best estimate of the pre-impact spectral
behavior of the comet comes from the day of impact, July 4th. We thus extracted the 5 ksec of
spectral data from the start of the July 4th observation just before the impact and compared it to a
spectrum of 20 ksec duration taken after impact and near the end of the July 4 th observation
(Figure 7). While we see the ~20% increase reported in the lightcurve photometry due to the
impact in summing the 300—1000 eV x-ray spectra from T1, with a 90% confidence range of
0.8 - 1.6 times the mean pre-impact count rate (equal to 2.7 x 10-13 ergs cm-2 s-1 or LX = 4.6 x
1014 ergs s-1), no significant changes are observed in the 500-700 eV region, and although the
post-impact spectrum is higher between 750-900 eV, the poor signal-to-noise makes it
impossible to draw any definite conclusions. These may be consistent with the observed solar
wind variations (Figure 5b), due to the passing of a shock flank by the comet on late July 4th
(Section IIIb).
V
Discussion and Conclusions
Here we summarize our findings observing Tempel 1 during the Chandra Deep Impact
observing campaign. First we list the morphological, temporal, and spectral results far in time
from the impact. This information gives us the comet’s behavior in the x-ray and sets the stage
for the July 4th Deep Impact experiment. We then discuss the changes observed due to the
impact, and compare the influence of the man-made cratering experiment with the two "natural"
outbursts found.
20
General Results

In terms of measured count rate, Tempel 1 is the second faintest comet ever detected in the x-ray. In terms
of 300 - 1000 eV x-ray luminosity, it ranks 4th faintest.

Tempel 1 in June - July 2005 is the first comet to be observed simultaneously and repeatedly by two
different x-ray instruments - the Chandra ACIS-S and the SWIFT HRT. Comparison of the measured
photometric fluxes and temporal trends for the two datasets shows good agreement.
 The standard CXE model works for this very well studied comet. The observed x-ray behavior (morphology,
lightcurve, and spectrum) for Tempel 1 can be explained with the collisionally thin model used for
2P/Encke 2003 (Lisse et al. 2005). The relatively normal line ratios of the Tempel 1 spectrum, when
compared to the unusual spectrum observed for 2P/Encke on 24 Nov 2003, allow us to determine that the
source of this unusual spectrum was the abnormal charge state of the solar wind ions incident at Encke
during the time of observation, not the low neutral number density in the comet's coma.
 The narrow June 30th flare coincides with detections of dramatically increasing outgassing rates by the Deep
Impact (Farnham et al. 2007) and Lowell Observatories (Schleicher et al. 2006), followed by an increased
solar wind flux. Roughly 50% of the increased emission during this flare was due to increases in the solar
wind flux incident on the comet; the rest is due to increasing cometary outgassing. While the outgassing
effect was difficult to detect, as jumps in observed x-ray emission rates from fluctuations in the incident
solar wind flux were large (i.e. much larger than those seen due to the material emitted due to the DI
experiment), long term monitoring of the comet over 75 ksec shows a definite trend attributable to an
outburst in neutral gas from the comet. The extent of the detectable x-ray emitting region increased by 50 ±
25%. These results represent the first demonstration of a direct link between rapid changes in a comet’s
Qgas(t) and the observed x-ray flux.
 The days-long, wide x-ray flare from the comet starting on July 8 is solar wind driven, and can be accounted
for entirely by changes in the solar wind time record. The July 8th x-ray flare did not increase the apparent
extent of the x-ray emitting region over the range seen during quiescent periods. Both naturally-occurring
x-ray flares from the comet were much stronger in intensity and total flux than that produced by the Deep
Impact experiment.
 The observed spectrum in and out of flare varies, while the observed morphology does not (except for the
June 30th flare, where a solar wind impulse coincided with a comet neutral gas outburst). The differences
in the spectrum can be explained by variation in the incident charge state ratio of the solar wind .
Deep Impact Experiment Results

No definitive x-ray flash was observed at impact. It is not clear if this is due to an intrinsic effect from the
impact, or the sensitivity of the Chandra/ACIS-S camera.

The overall emission increased by ~20% over the course of July 4th and stayed elevated through July 5.
There is no evidence for a huge change in the comet's x-ray emission behavior, moving the comet from
collisionally thin (to CXE) to collisionally thick behavior, as predicted by some pre-impact estimates (Lisse
et al. 2005, Schultz et al. 2005; Figure 8). Much of the reason for the small magnitude of the change is the
lack of production of a new active area, and the relatively small total amount of material, 10 6 - 107 kg,
liberated by the impact.

No large spectral shifts in the emission were seen due to the impact. The level of the observed changes is
within the range of statistical uncertainty.

Extended emission along the path taken by the excavated material was observed at the 5 level. The
detection is not strong enough to determine if the emission rate from the extension is changing with time, as
would be expected from an extended source of neutral gas, such as outgassing from ejected water ice. A
21
similar increase in the extent of x-ray emission was seen for the 'natural' outbursts of June 30th and July
8th.

VI
There is no evidence for a late, large upsurge in material emitted from the comet into the coma due to the
Deep Impact experiment, as reported by Willingale et al. 2006 based on the analysis of SWIFT data;
instead, the broad increase in x-ray flux seen in the combined Chandra/SWIFT x-ray lightcurves can be
explained by increases in the solar wind flux at the comet. The long term temporal behavior of the x-ray
luminosity, before, during, and after the Deep Impact experiment, can be matched using measures of the
coma neutral gas density (SWIFT, ESO CN, CALAR ALTO, SST, CARA) that show an elevated density
of material for the first two days after impact, then normal cometary behavior afterwards. There is also no
change in the emission morphology seen on July 8th, corresponding to that found on July 4th, during the
outgassing event reported by SWIFT.
Acknowledgements
We would like to thank V. Kharchenko, T. Sasseen, M. Torney, D. Willingale, and the SWIFT
project for many helpful discussions, as well as the anonymous reviewers whose comments
greatly improved the manuscript. The SOHO CELIAS heavy ion and MTOF proton monitor data
were
graciously
provided
by
the
University
of
Maryland
SOHO
project,
URL
http://umtof.umd.edu/. We are grateful for the cometary ephemerides of D.K. Yeomans et al.
found at URL http://ssd.jpl.nasa.gov/horizons.html used to reduce our data. Help with the MHD
modeling was provided by Y. Jia at the University of Michigan. C. Lisse was supported in part
by SMAO observing grants NAG GO45167X, and D. Bodewits and R. Hoekstra would like to
acknowledge support within the FOM-EURATOM association agreement.
VII
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26
VIII. Tables
Table 1.
Observation
Time
Chandra 9P/Tempel 1 Observation Times And Observing Parameters*
DOY
(UT 2005)

rh
Phase
(1027 Mol s-1)
(ks)
Total Counts in Energy Range
After Background Removal
500 –1000
(eV)
1610
181.7
1.51
0.875
40.6
245, 0.79
4.731
52
04 July 04:12
185.6
1.51
0.894
40.9
247, 0.34
5.011
64
300
1255
1560
248, 0.16
8.181
48
470
1170
1640
250, -0.14
6.311
28
570
1500
2070
2
08 July 09:13
189.6
1.51
1.51
0.902
41.1
0.916
41.2
(Degrees)
Texp
30 June 10:46
187.0
(Degrees)
QOH
300-500
(eV)
1130
05 July 17:12
(AU) (AU)
Heliocentric
Lon, Lat.
300-1000
(eV)
2740
10 July 23:39
191.2
1.51
0.931
41.4
252, -0.47
5.01
35
85
975
1060
13 July 15:31
194.8
1.51
0.945
41.5
253, -0.77
5.62
34
100
340
440
24 July 05:54
205.3
1.51
0.976
41.6
256, -1.36
5.62
34
175
815
990
*DOY is the center of the observation time. rh is the comet-Sun distance,  is the comet-Earth distance. The phase
angle is the angle Sun-comet-observer; the heliocentric ecliptic latitude and longitude are the comet’s position as
seen from the Sun, and Qgas is the gas production rate from the nucleus surface . 1 Schleicher et al. 2006. 2 Lisse et al.
2006. For reference, the heliocentric longitude and latitude of the Earth were in the range 281 - 292o, and 0o,
respectively.
Table 2. Solar Flares with Coronal Shock Waves During the Chandra/Tempel 1 Epoch
with Potential Interplanetary Consequences Monitored in Real Time *
FF#
Date
DOY
Time
Lat
Lon
Vs
x-ray
Opt.
TAU Classif. Classif.
606
03 June
154
23:55
S 17
E 10
499
100
C 6.2
607
16 June
167
20:10
N 09
W 87
989
20
XXXX
SF
same as above
608
01 July
182
5:01
N14
E 83
1200
15
C 5.3
XX
NARROW CME
609
07 July
188
16:29
N 09
E 03
1000
130
M 4.9
SN
HALO CME
610
08 July
189
6:34
N 19
E 31
330
100
B 5.5
SF
SOHO/LASCO NOT AVAILABLE
611
09 July
190
21:58
N 12
W 28
1651
100
M 2.8
1N
HALO CME
612
10 July
191
15:19
N 00
W 90
869
40
C 9.9
SF
SOHO/LASCO NOT AVAILABLE
613
12 July
193
16:27
N 09
W 64
1000
300
M 1.5
SF
PARTIAL HALO CME
614
13 July
194
15:27
N 11
W 90
1360
400
M 5.0
SF
HALO CME
615
14 July
195
11:20
N 11
W 90
1236
230
X 1.2
XX
same as above
616
16 July
197
7:13
S 10
W 74
667
200
C 2.2
XX
SOHO/LASCO NOT AVAILABLE
617
17 July
198
9:50
N 11
W130
2000
100
XXXX
XX
HALO CME
618
21 July
202
3:24
N 11
E167
2000
100
XXXX
XX
same as above
SN
Comments
SOHO/LASCO NOT AVAILABLE
619
24 July
205
13:36 N 11
E120 3000 100 XXXX XX
same as above
*FF# refers to the events that have been continuously "fearlessly forecasted" in real time since 1997 (see, for instance, Fry et al., 2003,
2004; Dryer et al., 2004). ‘Time’ refers to the start time of U.S. Air Force ground-based observations of metric Type II radio bursts
signifying a coronal shock. Lat and Long refer to the flare location on the Sun in solar polar coordinates. ‘Vs’ is the shock speed in the
low corona using the Type II's frequency drift rate. “Comments” refer to the time if-or when- coronagraph observations as reported by
NASA/SOHO operations scientists shortly after CME imagery is obtained. XXXX and XX indicate that flare x-ray and Optical (Halpha) classifications were not possible. Backside flare locations are based on time elapsed since West limb passage. TAU (in HMM)
refers to the shocks' piston-driving times assuming the flares' x-ray emission durations as a proxy. The HAFv.2 model uses only
Columns 2 and 4-8 as inputs.
27
Table 3. 0.3 - 1.0 keV Spectral Fits to the June - July 2005 Chandra ACIS-S Observations*
a
a
b
Line Fluxes (10-5 ph cm-2 s-1) and 90% C.L. (1.6)
July 04
July 05
July 08b
July 11
July 13
July 24
Line ID
…
…
CV f+r+i
8.0±6.0
<10.0
CVI Ly-
7.0±4.0
<6.0
5.2±4.0
NVI f+r+i
7.0±1.0
2.3±1.6
2.2±1.7
4.0±2.0
NVII Ly-
…
15.7±4.0
<1.0
<1.0
…
4.9±1.0
5.1±1.1
12.1±3.0
5.5±0.7
3.3±1.2
7.4±6.0
OVII r+i
5.7±1.4
2.2±0.5
1.9±0.7
12.6±4.0
3.5±1.0
1.9±0.8
2.5±1.0
OVIII Ly-
<1.0
1.5±1.1
1.0±0.5
1.0±0.6
5.3±3.5
1.4±0..8
1.0±0.7
1.1±0.7
OVII K-
774.6
1.0±0.5
1.2±0.7
0.9±0.3
1.2±0.4
<3.0
1.2±0.6
0.9±0.5
0.8±0.6
OVIII Ly-
849.1
0.5±0.2
1.2±0.5
0.8±0.2
1.0±0.2
3.4±1.0
1.5±0.2
< 0.5
0.5±0.3
OVIII Ly-
922.0
0.6±0.3
0.8±0.5
0.8±0.2
1.0±0.3
3.0±1.7
0.7±0.4
0.7±0.3
0.6±0.4
NeIX f+r+i
0.9
1.2
1.2
1.1
0.7
1.1
1.2
1.1
Eline
(eV)
June 30
307
358±60
508±100
112±40
174±60
490±190
…
367.5
18.7±5.6
33.4±7.4
8.8±4.6
3.9±6.0
41±19
<8.0
419.8
< 6.0
7.7±7.0
4.0±3.0
4.5±3.8
<25
500.3
3.2±1.8
6.8±3.0
3.0±1.0
<2.5
561.1
<0.5
…
<1.0
571.3d
4.8±1.1
24.2±3.0
653.5
1.6±0.7
712.7
Quiet state spectrum only.
June 30
b
Flare/ high state spectrum.
c
OVII f
2/dof
’<’ indicates a 90% confidence upper limit. dAverage of OVII 568.4 and 573.96 eV lines.
c
Table 4. Solar Wind Ion Abundances* Derived From the Spectral Fits in Table 3.
June 30a
June 30b
July 04
July 05
July 08a
July 11
July 13
July 24
S&C†
00
Bei 03
Kra 04
Otr 06
O8+
0.7 ± 0.28
0.4 ±0.09
1.0 ± 0.22
1.0 ± 0.3
0.7 ± 0.2
1.2 ± 0.2
1.4 ± 0.5
0.5 ± 0.5
0.35
0.13±0.03
0.13±0.05
0.35
C6+
3.0 ± 2.45
1.2 ±0.78
1.5 ± 1.82
0.8 ± 2.2
1.1 ± 1.8
0.7 ± 1.5
2.1 ± 2.9
0.6 ± 2.0
0.90±0.9
0.7±0.2
1.02
N7+
0.9 ± 0.4
0.4 ±0.13
0.8 ± 0.3
0.4 ± 0.3
0.3 ± 0.4
0.5 ± 0.3
0.9 ± 0.6
0.7 ± 0.5
0.06±0.02
…
0.03
Date
Ion
1.59
0.03
* - Solar wind abundances (relative to O7+) derived from spectra ± 90% C.L. † - Slow wind model parameters from Schwadron and Cravens (2000); Bei 03 =
Beiersdorfer et al. (2003); Kra 04 = Krasnopolsky et al.. (2004); Otr 06 = Otranto et al. (2006).
28
IX. Figures
Figure 1. ACE/SWICS solar wind parameters. The top panel shows densities and relative compositions of a few
solar wind heavy ions, as indicated by the legend at the right. The middle panel indicates the velocity of some of
these ions, indicating a very tight correspondence of all speeds. The labels refer to the text. The lowest panel shows
thermal speeds, which also appear to track each other very well. Unusual states of the solar wind can be easily
located by searching for deviations from normal in the O7+/O6+ (brown/olive curve) in the top panel. DOY 178
corresponds to June 28, DOY 213 to August 1.
29
Figure 2. Solar wind interplanetary
magnetic field structures, estimated
using the HAFv2 model at the times
of the Chandra observations. Red :
magnetic polarity away from Sun;
blue : toward the Sun. The
coordinate system is semi-inertial
(HEQ, heliocentric solar equatorial)
with 0o longitude at the intersection
of the solar equatorial plane and the
ecliptic plane. (The ecliptic plane is
inclined by 7.25o to the solar
equatorial plane. The HAF model
resolution is 5o in heliolatitude and
heliolongitude, so the ecliptic plane
and equatorial plane IMF plots are
quite similar, particularly in areas
near the intersection of the equatorial
and ecliptic plane. Since T1 is near
the 0-180 degree line on the IMF
lots, its low heliocentric ecliptic
latitude puts it close to the solar
equatorial plane during this period.)
There are no obvious ICME
structures in the
comet-ward
directions during the times of the
Chandra-Tempel 1 observations
except for the west flank of the
interplanetary shock on July 4 (DOY
185). Interplanetary shocks from
parent solar flares (indicated in Table
2 as FF#611, #613, #614, #615, and
#617) passed ACE, respectively on
12, 16, 16, 17, and 19 July as given
by real-time-reported plasma and
magnetometer Level 1 data from
ACE (Zwickl et al. 1998.) Their
eastern
flanks
however
are
considered to have been very weak at
Tempel 1. On the other hand, CIRs
are indicated to intercept the comet
on June 30 (DOY 181), July 8 (DOY
189), and July 24 (DOY 205). On
the basis of the model’s output with
ACE proton data, the simulated
FF#611 shock (see the panels for
July 11/DoY 192 and 13/DoY 194)
arrived a day earlier than the actual
one; thus, the implied eastern flank’s
shock speed and strength at the
comet are too high.
30
Figure 3. Chandra ACIS-S3 images of the 9P/Tempel 1 x-ray emission. Images were created from all photons
with 0.5 < Energy < 0.7 keV, smoothed with a 5x5 pixel Gaussian filter. Note that the imagery from the July 24th
observation has been omitted for ease of presentation; the observed structure is nearly identical to that found on July
14th. North is up, East to the left, and the projected direction to the Sun is to the right in each image. The images
cover an extent that is roughly 3.5 x 10 5 km by 3.5 x 105 km. Inset : July 4th KPNO 4m R-band image, rotated and
scaled to match the Chandra images (Farnham et al. 2007). Arrows denote the path of the ejecta, which is moving
down and to the right (SW). The position of the nucleus in the KPNO and Chandra images is denoted by a black
cross. a) The brightness in each image has been autoscaled to the maximum value of the image, enhancing the
extension in the fainter, non-flare observations. b) Same as a), but with each image hard-scaled to the maximum
brightness on the July 11th image.
31
Lightcurve
Proton flux
Gas production
Figure 4. a) Total 0.3 to 1.0 keV Chandra count rate versus time for the period June 30 - July 14, 2005, with
500 second bins smoothed by an 11 point boxcar filter (small solid squares). The sharp impulse on DOY 181.7 (June
30) is in strong contrast to the wide observed increase in x-ray count rate in DOY 189.6 (July 8). Points within a
boxcar width of the beginning or end of each Chandra visit to the comet are not shown. The dashed line marks the
time of impact of the DI spacecraft. Also plotted, on an arbitrary scale, are the interpolated and scaled Q gas temporal
trends (triangles), np*vp solar wind trends (stars), SWIFT photometry (Willingale et al. 2006; open large squares)
and np*vp*Qgas trends (solid line). Qgas has been estimated using the QOH outgassing rates from Schleicher et al.
(2006), shifted by -1 day, combined with the 16 µm Spitzer lightcurve data of Lisse et al. (2006).
32
Figure 4. b) Chandra lightcurve phased to the diurnal rotation rate of 1.6 days. Error bars
are 1. Zero phase corresponds to Day 181 00:00 UTC of Year 2005.
33
Figure 5a. Spectra (background subtracted) for a sample of the Chandra ACIS observations plotted from 300
to 1000 eV. Each spectrum is offset from the previous by 0.2 count s -1 keV-1. 1 error bars are shown. The spectra
show remarkable day-to-day changes in the emission between 500 and 650 eV, attributable to changes in the OVII
and OVIII emission (see text). The natural x-ray outbursts of June 30th and July 8th caused much larger changes
than the man-made outburst of July 4-5. The spectra of June 30th (flare) and July 8th (quiet) are very similar, while
the July 8th (flare) spectrum, though noisy, has much stronger C,N, and O charge exchange emission lines,
suggesting the dominant driver for the observed spectral changes is the solar wind ion composition.
34
Figure 5b. Ion abundances for O8+ (top) and C6+ (middle) relative to O7+. In the lower panel, the MHD results
for the proton velocity are given for comparison. In all panels, the lines are ACE/SWICS data, as obtained with the
'simple extrapolation' (corotational+distance), and the error bars are 90% C.L. (±1.6. The dashed lines are
averaged slow wind abundances from Schwadron & Cravens (2000).
35
Figure 6. Neutral coma gas model used for the spectral modeling. Line styles are: solid black – water; dashed
black – CO; dotted black – H; solid grey – OH; dotted grey – C; dash-dotted grey – O. QH2O = 1 x 1028 mol/s, H2O =
1.3 x 105s; QCO = 0.12 * QH2O, CO = 9.0 x 105s; QCO2 = 0.05 * QH2O, CO2 = 3.3 x 105s. Lifetimes are from Huebner
et al. (1992) for an active Sun, scaled to 1.51 AU. Note that the neutral abundance is dominated by water and its
dissociation products H2O, OH, and O throughout the coma.
36
Figure 7. Comparison of the Pre- and Post- Impact spectra from 04 July. Shown are the spectra (with ±1 error
bars) with pre-impact in open squares and post-impact in filled circles, over plotted with the best fitting multiemission line CXE models (dashed for pre-impact and solid line for post-impact). The post-impact spectrum has
been scaled to the pre-impact levels; an overall change in the 300 - 1000 eV total flux of 18 ± 5 (1) % was found.
Marginal differences in the pre- and post-spectra are seen, if at all, in the 800-900 eV region.
37