# Peculiar temporal and spectral features in highly obscured HMXB pulsar IGR
J16320-4751 using

Varun,1 Nirmal Iyer,2,3 Biswajit Paul4
1Aryabhatta Research Institute of Observational Sciences, Nainital, India
263001
2Department of Physics, KTH Royal Institute of Technology, Stockholm, Sweden
10691
3The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova University Centre,
Stockholm, Sweden 10691
4Raman Research Institute, Bangalore, India 560080
E-mail<EMAIL_ADDRESS>
###### Abstract

IGR J16320-4751 is a highly obscured HMXB source containing a very slow
neutron star ($P_{spin}\sim 1300$ sec) orbiting its supergiant companion star
with a period of $\sim$9 days. It shows high column density ($N_{H}\sim
2-5\times 10^{23}$ $cm^{-2}$) in the spectrum, and a large variation in flux
along the orbit despite not being an eclipsing source. We report on some
peculiar timing and spectral features from archival observation of this source
including 8 observations taken during a single orbit. The pulsar shows large
timing variability in terms of average count rate from different observations,
flaring activity, sudden changes in count rate, cessation of pulsation, and
variable pulse profile even from observations taken a few days apart. We note
that IGR J16320-4751 is among a small number of sources for which this
temporary cessation of pulsation in the light curve has been observed. A time-
resolved spectral analysis around the segment of missing pulse shows that
variable absorption is deriving such a behavior in this source. Energy
resolved pulse profiles in 6.2-6.6 keV band which has a partial contribution
from Fe Kα photons, show strong pulsation. However, a more systematic analysis
reveals a flat pulse profile from the contribution of Fe Kα photons in this
band implying a symmetric distribution for the material responsible for this
emission. Soft excess emission below 3 keV is seen in 6 out of 11 spectra of
observations.

###### keywords:

binaries: general — pulsars: individual (IGR J16320-4751) — X-rays: binaries —
methods: data analysis

##  1 Introduction

The study of high mass X-ray binaries (HMXB) gives us insight into many areas
of modern astrophysics including accretion physics, stellar wind, stellar
evolution, and formation of double compact binaries. HMXBs were among the
first astrophysical sources studied using space based X-ray telescopes. All
sky surveys conducted before 2000 discovered 110 HMXBs (Liu et al.,, 2000)
using soft X-ray instruments. A majority of these sources were Be binaries and
only a few systems were classified as supergiant binaries. However, there was
a systematic bias in this sample due to the operating energy range ($<$12 keV)
of these surveys. The sample of HMXBs has been updated in the past two decades
with hard X-rays ($>$17 keV) surveys conducted by INTEGRAL and Swift
observatory. In particular, the number of supergiant HMXBs increased by almost
a factor of three (Walter et al.,, 2015) and a large number of these were
characterized as highly obscured HMXBs. This previously undetected subclass of
HMXB shows large intrinsic photoabsorption ($N_{H}>10^{23}$ $cm^{-2}$) in
their spectra (Zurita et al.,, 2006; Rodriguez et al.,, 2006). IGR J16318-4848
is the most extreme case in this subclass with the highest known intrinsic
absorption ($N_{H}\sim 10^{24}$ $cm^{-2}$) (Ibarra et al.,, 2007). Most of the
highly obscured sources have short orbital periods and possibly undergoing a
transition to Roche lobe overflow (Walter et al.,, 2015). Highly obscured
sources IGR J16320-4751 and EXO 1722-363 have larger orbital period ($\sim$10
days) and hence are unlikely to be in the Roche lobe transitioning phase.

IGR J16320-4751 was discovered in February 2003 by IBIS/ISGRI detector (Lebrun
et al.,, 2001) during a target of opportunity (ToO) observation of another
binary source 4U 1630-47 (Tomsick et al.,, 2003). It was revealed later that
this source was associated with ASCA source AXJ1631.9-4751, observed in 1994
(Sugizaki et al.,, 2001). Its X-ray position was determined accurately at RA =
$16^{h}32^{m}01^{s}.9$, Dec = $-47^{\circ}52\arcmin 27\arcsec$ using
observations (Rodriguez et al., 2003b, ). observations also revealed timing
variability in the source and X-ray pulsation of $P=1309\pm 40$ s associated
with the spin period of the neutron star (NS) in the system (Lutovinov et
al.,, 2005). This binary is located at a distance of $\sim$3.5 kpc (Coleiro et
al.,, 2013). A strong modulation at a period of $8.96\pm 0.01$ days in the
long term Swift-BAT corresponds to the orbital period of NS around its
companion star (Corbet et al.,, 2005). The presence of several emission lines
in the near infra-red (NIR) spectrum (e.g. Pa(7-3), Br(17-4), and HeI(7-4))
indicate that the companion star might be an O/B supergiant (Chaty et al.,,
2008). Companion star is not visible in any optical bands corresponding to the
X-ray position of the source due to the high intrinsic absorption.

X-ray light curves of IGR J16320-4751 show large variability at various time
scales, from seconds-minutes (Rodriguez et al., 2003a, ) to days-months
(Foschini et al.,, 2004). The pulse profile and fraction are also seen to vary
with time. However, the pulse fraction is found to be constant with energy in
both and INTEGRAL energy bands (Rodriguez et al.,, 2006). The X-ray spectrum
of this source is similar to other highly absorbed HMXBs, consisting of a
power-law, prominent iron emission lines along with a large absorption column
($N_{H}\simeq 2-5\times 10^{23}cm^{-2}$). Despite this large column density,
the presence of an excess of soft photons below $\sim$3 keV has been noted
(Rodriguez et al.,, 2006). The absorption, $N_{H}$, and line emission (EW of
Fe $K_{\alpha}$) are correlated with each other, as expected from a system
with the absorbing material also yielding fluorescence lines (García et al.,,
2018).

Attempts have been made before to understand the orbital variations of IGR
J16320-4751 by proposing for and studying multiple observations along a single
orbit (Heras et al.,, 2009; García et al.,, 2018). In this work, we use these
observations to study the variable pulse profile of the neutron star. Iron
emission lines have been found in spectra of a large number of HMXBs and have
proven to be a useful tool to study stellar wind properties in these systems
(Torrejón et al.,, 2010; Pradhan et al.,, 2018; Aftab and Paul,, 2017). We
study the pulse profile in the energy range of these iron line emissions to
look for clues about the overall pulsation behavior. From these studies, we
gather evidence for some peculiar temporal and spectral features in this
binary system. These studies are detailed in the following sections. We give
details of the instruments and observations used in § 2. Details on the timing
and spectral analysis undertaken are given in § 3. Finally, we present our
results followed by a discussion of these findings in § 4.

##  2 Observation and data reduction

We used archival data from observations of IGR J16320-4751 for this work. is a
space-based high energy mission launched by the European Space Agency (ESA)).
A set of three European Photon Imaging Cameras (EPIC) serve as the primary
instruments on board (Jansen et al.,, 2001). The EPIC cameras consist of a
single PN CCD camera and two MOS (Metal Oxide Semi-Conductor) CCD cameras. The
PN camera is made of twelve CCD units embedded into a single silicon chip. The
EPIC PN camera has a moderate effective area ($\sim$1000 $cm^{2}$) and good
spectral resolution ($\Delta E/E\sim 20-50$) in its nominal energy range
(0.15-15 keV). The observations used for this paper are done in the full-frame
imaging mode, in which EPIC PN has a time resolution of 73.4 ms. IGR
J16320-4751 was observed 11 times with between March 2004 and September 2008.
Details of these observations are given in Table 1. The observation numbers
given in Table 1 are used to identify individual observations in the remainder
of this paper. 8 of 11 observations (Obs #3 - Obs #10) used in this work were
carried out during 14-26 August 2008 at different orbital phases of the same
binary orbit. This allows us to investigate orbital variations in detail over
the same orbit.

Table 1: Table Of Observations

Obs.No. | Obs. Id | Date-Start | Tstart | Date-End | Tend | Duration[s] | Orbit Phase
---|---|---|---|---|---|---|---
1 | 0128531101 | 04-03-2004 | 20:58:26 | 05-03-2004 | 03:12:27 | 22441 | 0.769-0.797
2 | 0201700301 | 19-08-2004 | 13:28:23 | 20-08-2004 | 03:20:40 | 49337 | 0.419-0.483
3 | 0556140101 | 14-08-2008 | 22:44:10 | 15-08-2008 | 01:14:43 | 9033 | 0.401-0.413
4 | 0556140201 | 16-08-2008 | 17:41:23 | 16-08-2008 | 19:55:16 | 8033 | 0.601-0.611
5 | 0556140301 | 18-08-2008 | 13:36:12 | 18-08-2008 | 15:33:26 | 7033 | 0.804-0.814
6 | 0556140401 | 20-08-2008 | 07:36:07 | 20-08-2008 | 10:43:20 | 11233 | 0.999-0.013
7 | 0556140501 | 21-08-2008 | 07:04:12 | 21-08-2008 | 07:41:30 | 2238 | 0.108-0.110
8 | 0556140601 | 22-08-2008 | 03:55:58 | 22-08-2008 | 07:22:36 | 12397 | 0.204-0.220
9 | 0556140701 | 24-08-2008 | 18:30:00 | 24-08-2008 | 21:00:33 | 9033 | 0.494-0.506
10 | 0556140801 | 26-08-2008 | 13:34:33 | 26-08-2008 | 16:15:06 | 9632 | 0.694-0.706
11 | 0556141001 | 17-09-2008 | 01:25:41 | 17-09-2008 | 03:31:14 | 7555 | 0.084-0.094

  * 1

Notes: Orbital phases of different observations has been calculated using
orbital period = $8.991\pm 0.001$ days (obtained in this work) and taking
reference epoch at the mid point of observation #9 (54702.82310 MJD).

To analyze the timing and spectral properties of IGR J16320-4751 from the EPIC
PN event files, we extracted source information using a circular region of
radius $40\arcsec$ centered on the source position. The analysis was done
using XMM Science Analysis System (SAS version 20.0.0). Background events were
extracted from a similar source-free circular region. Both light curves and
spectra were extracted from these regions and the background was subtracted
from the source counts. The light curves were also barycenter corrected using
the SAS tool barycen. The spectra were grouped to have a minimum of 25 counts
in each bin with specgroup and fit using XSPEC V12.10.1.

Figure 1: _Left panel:_ Orbital profile of IGR J16320-4751 using long term
Swift-BAT data. Two cycles of orbital profile are shown on which orbital
phases of all observations are marked in the first cycle. _Right panel:_
Actual Swift-BAT light curve for August 2008 during which 8 observations were
taken at different phases covering the entire orbit. The vertical dotted line
in this panel corresponds to orbital phase 0.

We used the Swift-BAT long term light curve from February 2005 to December
2020 and searched for periodicity to determine a period of $8.991\pm 0.001$
days. The error on the spin period value is found bootstrap sample of 1000
light curves (Filippova et al.,, 2004; Boldin et al.,, 2013). We generated
copies of original light curves using a Gaussian random number generator
taking each data point with value and error on it as mean and standard
deviation respectively. The best period was determined from each light curve
using the tool efsearch. Finally, we computed the standard deviation of the
distribution of period values. This value is similar to the orbital period
reported in (García et al.,, 2018), with a lower error owing to a longer
dataspan. The orbital profile created with this period folded using a
reference epoch at 54702.82310 MJD is shown in the left panel of Figure 1. It
is interesting to note that although the source is not an eclipsing HMXB
system, it shows a very strong orbital intensity modulation. The observations
we use in this paper, are marked on the orbital profile in Fig. 1 at their
respective orbital phase values. The actual Swift-BAT light curve during the
single binary orbit covering 8 observations in August 2008 is shown on the
right panel of Figure 1.

##  3 Analysis and Results

###  3.1 Timing

Pulsations are immediately apparent in the EPIC PN light curves of IGR
J16320-4751. Before analyzing these pulsations, we inspected the light curves
to see if any other significant features were visible. A noticeable feature is
a difference in the count rate seen in the light curves from different
observations. Except for observation #1 which shows a steady non-variant light
curve, all other observations display variability of some kind. Flaring
activity for a few ks can be seen in observations #2 and #4 during which the
count rate increases by a factor of $\sim$5 compared to the average count rate
in the light curve (see Figure 2). Observations #10 and #11 show comparatively
shorter flares (2 and 1 respectively) for a few 100 seconds. The count rate in
observation #3 is steady for 8 ks and it starts to fall in the end. The count
rate in observations #5 and #6 show gradual increase whereas observations #8
and #9 shows random rise and fall.

Figure 2: Light curves from all observations in Table 1 are plotted with 20.0
sec time bins. Time axis has been rescaled to start from zero for all
observations.

We observed that flares affect the shape of the pulse profile and the results
of our pulse profile analysis. So we removed these sections from light curves
before creating pulse profiles. We searched for periodicity in each
observation using the tool efsearch. This gave us a most probable period value
of 1305.9 seconds, which was obtained by using data combined from three
observations made on two consecutive days, viz. obs #6, #7, and #8. We created
a template pulse profile by combining light curves from three observations and
pulse profiles from each observation in energy band 0.2-10.0 keV using the
most probable period and a reference epoch at MJD 54692.0, with the tool
efold. All pulse profiles (created with 32 phasebins) are shown in Figure 3.
We note that the pulse profiles show a broad single peaked feature. The pulse
profile of obs #8 also shows a small inter-pulse. Given that 8 of 11
observations were made during a single orbit, we attempted to use the pulse
profiles to determine the orbital characteristics of the pulsar using phase
lags between observations. The template profile is correlated against the
profiles from individual observations to find the phase lag between
observations. As seen from the pulse profiles in Figure 3, there is a
significant shift in the position of the primary pulse peak (for example
between obs #5 and obs #8) which cannot be easily explained in terms of
orbital changes alone. For the orbital motion to lead to such large phase-lags
($\sim$0.4), would require an extremely large mass function
$\sim$280$M_{\odot}$. Apart from orbital motion, a change in pulsar period can
also lead to shift in the phases of pulse profiles if all the light curves are
folded with a constant period. The estimated period derivative to derive a
phase shift of $\sim$0.4 comes out to be 5$\times 10^{-4}$ ss-1 which is much
larger than the period derivatives of 10-12 ss-1 seen for HMXBs (Ho et al.,,
2014). It is more likely that these phase-lags are caused by other reasons as
elaborated in the discussion section.

Figure 3: Pulse profiles obtained from different observations. Top left panel
shows a pulse profile created from three close observations taken. Later pulse
profiles are for different individual observations. Each pulse profile
contains a marker for the observation number as presented in the Table 1. Two
dashed horizontal lines are marked in each pulse profile at 0.9 and 1.1 to
demonstrate the variability among them.

We noted a curious absence of pulsations in obs #11. The pulsations are seen
as distinct $\sim$1300 s modulations in all light curves except obs #11. We
also noticed many instances of pulse to pulse variations in other observations
where the characteristic pulse is either missing or significantly different
from the broad single peak. We investigated this in more detail by overlaying
the pulse profile created from the observation over the entire duration of the
light curve and found pulse-to-pulse variation in 6 out of 11 observations of
IGR J16320-4751. There are several instances of single missing pulses as well
as multiple missing pulses together. In Figure 4, we show the light curve of
observation #2 of IGR J16320-4751. Missing pulse intervals are marked in the
light curves as “MP” (Left panel). An expanded view near 30 ks in this
observation is shown to display this more clearly where the first interval has
clear pulses and two pulses are missing in the second interval (Right panel).
A comparison of spectral parameters between these two intervals is given in
the spectroscopy section.

Figure 4: _Left Panel:_ Light curve from observation #2 of IGR J16320-4751
using 20 s binning time. Time intervals during which pulses are missing are
marked as MP._Right Panel:_ An expanded view near 30 ks in this observation.
First interval has clear pulsation whereas two pulses are missing continuously
in the second interval.

Figure 5: Energy resolved pulse profiles from observation #6 (Left) and #9
(Right). Top three panels show pulse profiles from low energy band (3.0-6.0
keV), mid energy band (6.2-6.6 keV), and high energy band (7.2-10.0 keV)
respectively.

The X-ray spectrum of this source showed strong Fe emission lines as discussed
in the next section. To check the timing properties of this emission feature,
we extracted pulse profiles in three different energy bands, viz. the low
energy band (3.0-6.0 keV), the mid-energy band around Fe $K_{\alpha}$ line
emission (6.2-6.6 keV), and the high energy band (7.2-10.0 keV). Such energy-
resolved pulse profiles are shown in Figure 5 for observations #6 and #9. As
seen in the figure, there is a significant amount of pulsation seen in the mid
energy band. Mid energy band pulse profile has a contribution from the Fe
$K_{\alpha}$ photons and the continuum photons in that energy band. We
calculated the pulse profile of only the Fe $K_{\alpha}$ photons by removing
the contribution of continuum photons. Since the continuum pulse profile in
the band below and above the line energy are nearly identical, we assumed the
pulse profile of continuum photons in the 6.2-6.6 keV energy band to be the
average of the two other profiles. We subtracted the same from 6.2-6.6 keV
energy band pulse profile to obtain the pulse profile of the Fe $K_{\alpha}$
photons. We assumed that the mid energy profile is the weighted average of Fe
$K_{\alpha}$ photons pulse profile and continuum photons pulse profile,
weights being equal to the proportion of flux they contribute to this energy
band. The normalized pulse profile in the three energy bands, as shown in
Figure 5 have been used for the same.

$I_{Avg}(n)$ is calculated from the normalized pulse of the low and high
energy pulse profile as :

$\displaystyle I_{Avg}(n)=\frac{I_{Low}(n)+I_{High}(n)}{2}$ (1)

Here we have assumed that pulsation behavior of continuum in whole band is
same. Intensity in $n^{th}$ bin due to Fe $K_{\alpha}$ line photons
$I_{Line}(n)$ is given as:

$\displaystyle
I_{Line}(n)=1+\frac{BW+EQW}{EQW}\Bigg{(}\Big{(}I_{Mid}(n)-1\Big{)}-\Big{(}1-\frac{EQW}{BW+EQW}\Big{)}\Big{(}I_{Avg}(n)-1\Big{)}\Bigg{)}$
(2)

In this equation EQW is Equivalent Width of Fe $K_{\alpha}$ line emission, BW
is total band width of mid energy range. These Fe $K_{\alpha}$ photons pulse
profiles are plotted in bottom panels in Figure 5. We do not observe any
pulsed behavior by Fe $K_{\alpha}$ photons.

###  3.2 Spectroscopy

The continuum spectrum of IGR J16320-4751 can be modeled either with a
phenomenological powerlaw model (Pradhan et al.,, 2018) or with a thermally
comptonized model (COMPTT, Titarchuk, (1994)) to interpret the data in terms
of physical conditions close to the surface of NS (García et al.,, 2018). In
addition, a powerlaw with high energy cutoff model has been found to fit the
spectral data as well (Rodriguez et al.,, 2006; García et al.,, 2018). We
choose this model to describe the continuum in 3.0-11.0 keV energy range
because it gives a better value of test statistic $\chi^{2}$ compared to the
simple powerlaw model. We used 3 Gaussian components for line-like features in
6.0-7.5 keV energy range in a similar approach followed by García et al.,
(2018) by freezing their width to zero. The use of 3 Gaussian components was
chosen compared to 2 Gaussian components as it improves the residuals. Only
one Gaussian component was needed in observation 1 fit and two Gaussian
components were required for 4 observations (see Table 3). While this
describes the spectral data well at energies $>$3 keV for all the observations
(shown in the left panel of Figure 6), at lower energies a soft excess is seen
(right panel of Figure 6) in the some of the observations. To model this soft
excess we tried three components viz: 1. additional absorption with partial
covering 2. unabsorbed bremsstrahlung and 3. absorbed black-body. Neither of
these models could account for the soft excess completely. A partial covering
absorption component either gives unacceptable $\chi^{2}$ for some
observations or large residuals for other observations. The bremsstrahlung
model results in unreasonably high temperature ($\sim$200 keV) for the
electron plasma. The blackbody model parameters are not well constrained when
describing the soft excess data. To quantify the flux of the soft excess
component, we used the unabsorbed blackbody model and notice the value of the
norm of this spectral component and convert it into flux. We detect the
presence of soft excess in 6 out of 11 observations (see Table 2). The flux of
this emission is in the range of (0.4-2.4)$\times 10^{-13}$ erg cm-2 s-1. For
other observations, this soft excess flux is consistent with zero.

Figure 6: Spectrum of IGR J16320-4751 with a model fit from observation 8.
_Left panel:_ Spectrum in range 3.0-11.0 keV. _Right panel:_ Spectrum in range
0.3-11.0 keV. Bottom panels shows residuals of fitted model. Table 2: Table of
parameter values for the continuum component of spectral model.

Obs. No. | Parameters
---|---
| $N_{H}$ |  $E_{cut}$(keV) |  $E_{fold}$(keV) | $\Gamma$ | Norm2 | Soft Excess | $\chi^{2}_{red}$
1 | $19^{+3}_{-6}$ | $8^{+1}_{-2}$ | $7^{+8}_{-1}$ | $0.8^{+0.3}_{-0.9}$ | $3^{+2}_{-2}$ | - | 1.19
2 | $15.0^{+0.3}_{-0.9}$ | $7.4^{+0.2}_{-0.3}$ | $15^{+2}_{-1}$ | $0.43^{+0.05}_{-0.14}$ | $3.4^{+0.3}_{-0.7}$ | $0.9^{+1.1}_{-0.4}$ | 1.21
3 | $27^{+1}_{-3}$ | $8.1^{+0.4}_{-0.6}$ | $11^{+4}_{-3}$ | $0.5^{+0.1}_{-0.1}$ | $4.6^{+0.1}_{-0.9}$ | $0.4^{+1.4}_{-0.2}$ | 1.32
4 | $19.9^{+0.5}_{-1.0}$ | $8^{+1}_{-1}$ | $29^{+16}_{-9}$ | $0.39^{+0.09}_{-0.38}$ | $7^{+1}_{-1}$ | $0.6^{+1.3}_{-0.2}$ | 1.13
5 | $19^{+1}_{-1}$ | $7.0^{+0.9}_{-0.6}$ | $18^{+6}_{-5}$ | $0.3^{+0.2}_{-0.2}$ | $5^{+2}_{-2}$ | - | 1.17
6 | $22^{+3}_{-2}$ | $7.8^{+0.4}_{-0.5}$ | $14^{+5}_{-3}$ | $0.43^{+0.09}_{-0.11}$ | $4.6^{+0.4}_{-0.9}$ | - | 0.94
7 | $21^{+1}_{-1}$ | $7.6^{+0.6}_{-0.7}$ | $9^{+4}_{-3}$ | 0.43 | $3.8^{+0.2}_{-0.2}$ | - | 0.65
8 | $20.4^{+0.5}_{-0.3}$ | $7.7^{+0.2}_{-0.6}$ | $16^{+2}_{-2}$ | $0.24^{+0.06}_{-0.12}$ | $6.3^{+0.8}_{-1.3}$ | $2.4^{+3.4}_{-0.1}$ | 0.92
9 | $46^{+1}_{-1}$ | $8.7^{+0.4}_{-0.5}$ | $14^{+6}_{-5}$ | 0.24 | $3.6^{+0.1}_{-0.1}$ | - | 1.16
10 | $20^{+1}_{-2}$ | $7.2^{+0.9}_{-0.6}$ | $10^{+3}_{-2}$ | $0.3^{+0.2}_{-0.2}$ | $4^{+2}_{-1}$ | $1.0^{+3.6}_{-0.6}$ | 0.91
11 | $18^{+1}_{-4}$ | $7.2^{+0.3}_{-0.7}$ | $16^{+2}_{-1}$ | $0.3^{+0.1}_{-0.5}$ | $4.0^{+2.0}_{-0.9}$ | $0.4^{+0.3}_{-0.1}$ | 1.13

  * •

Notes: Equivalent hydrogen column $N_{H}$ is in units of $10^{22}$ atoms cm-2.
The normalization of photon power law is in units of photons $10^{-3}$ keV-1
cm-2 s-1 at 1 keV. The soft excess in 0.3-3.0 keV energy range is given in the
units of $10^{-13}$ erg cm-2 s-1.

Table 3: Table of parameter values for the line components of spectral model.

Obs. No. | Parameters
---|---
|  $E_{L1}$ (keV) | $Norm_{L1}$ |  $E_{L2}$ (keV) | $Norm_{L2}$ |  $E_{L3}$ (keV) | $Norm_{L3}$
1 | $6.38^{+0.09}_{-0.11}$ | $0.3^{+0.2}_{-0.2}$ | - | - | - | -
2 | $6.398^{+0.002}_{-0.007}$ | $1.78^{+0.09}_{-0.10}$ | $6.68^{+0.07}_{-0.07}$ | $1.41^{+0.8}_{-0.8}$ | $7.04^{+0.03}_{-0.03}$ | $2.7^{+0.8}_{-0.7}$
3 | $6.414^{+0.018}_{-0.004}$ | $4.5^{+0.3}_{-0.3}$ | - | - | $7.06^{+0.03}_{-0.04}$ | $7^{+2}_{-2}$
4 | $6.414^{+0.005}_{-0.011}$ | $6.8^{+0.4}_{-0.4}$ | $6.75^{+0.10}_{-0.08}$ | $5^{+3}_{-3}$ | $7.05^{+0.04}_{-0.04}$ | $12^{+3}_{-3}$
5 | $6.430^{+0.025}_{-0.006}$ | $5.2^{+0.4}_{-0.4}$ | - | - | $7.05^{+0.05}_{-0.06}$ | $7^{+3}_{-3}$
6 | $6.412^{+0.007}_{-0.023}$ | $3.1^{+0.3}_{-0.6}$ | $6.62^{+0.08}_{-0.23}$ | $4^{+2}_{-2}$ | $7.02^{+0.05}_{-0.04}$ | $6^{+2}_{-3}$
7 | $6.40^{+0.02}_{-0.02}$ | $2.4^{+0.5}_{-0.5}$ | - | - | $7.0^{+0.1}_{-0.2}$ | $5^{+4.0}_{-4}$
8 | $6.418^{+0.005}_{-0.016}$ | $8.5^{+0.3}_{-0.3}$ | $6.73^{+0.02}_{-0.04}$ | $11^{+3}_{-3}$ | $7.036^{+0.028}_{-0.007}$ | $26^{+3}_{-3}$
9 | $6.418^{+0.006}_{-0.002}$ | $10.0^{+0.5}_{-0.5}$ | $6.73^{+0.04}_{-0.06}$ | $10^{+3}_{-3}$ | $7.06^{+0.02}_{-0.03}$ | $19^{+4}_{-2}$
10 | $6.410^{+0.007}_{-0.004}$ | $3.8^{+0.4}_{-0.3}$ | - | - | $7.00^{+0.06}_{-0.05}$ | $7^{+3}_{-3}$
11 | $6.42^{+0.01}_{-0.04}$ | $3.9^{+0.4}_{-0.3}$ | $6.66^{+0.08}_{-0.06}$ | $4^{+4}_{-3}$ | $7.03^{+0.02}_{-0.03}$ | $9^{+3}_{-3}$

  * •

Notes: The normalisation of first Gaussian is given in the units of $10^{-4}$
photons cm-2 s-1. The normalisation of second and third Gaussian’s is given in
the units of $10^{-5}$ photons cm-2 s-1.

The behavior of Fe $K_{\alpha}$ photon with pulse phase can also be probed
using pulse phase-resolved spectroscopy. We checked for variation of the flux
of the Fe $K_{\alpha}$ line in the phase-resolved spectra of two observations:
one with the largest exposure time (observation #2) and another showing the
largest equivalent width of the Fe $K_{\alpha}$ line in phase averaged
spectrum (observation #8). We extracted ten phase-resolved spectra with a
phase bin width of 0.1 from these observations. We fit all the phase-resolved
spectra with the same model as used for the phase averaged spectrum. Line
energies for three Gaussian’s were frozen at phase averaged value and line
widths were fixed to zero for all the phase averaged spectra. We noted the
norm of the Gaussian corresponding to the Fe $K_{\alpha}$ line from all the
phase resolved spectra. We did not find any statistically significant
variation in the value of this spectral parameter with phase for both the
observations (See the left panel of Figure 7). Hence, this additional check
also validates that iron line photons do not show any variation with the pulse
phase. We discuss the implications of such behavior on the distribution of
matter.

Figure 7: _Left panel:_ The behavior of normalizaion of Fe $K_{\alpha}$ line
in the spectra of IGR J16320-4751 with pulse phase. _Right panel:_ Spectra of
IGR J16320-4751 from two adjescent intervals in the light curve of observation
#2. Solid line represents the model of the first interval rescaled to the flux
level of the second interval spectra.

We performed a time-resolved spectral analysis to understand the physical
reason behind the peculiar pulse to pulse variations in the light curves of
IGR J16320-4751. We selected an interval in observation #2 where the pulsation
from the source disappears for two consecutive cycles and another interval of
the same duration just before this interval where the pulsation is seen
clearly. We accumulated the phase-averaged source and background spectra in
3.0-11.0 keV from these intervals. In the right panel of Figure 7, we show the
spectra from these two intervals. We fitted both these spectra simultaneously
with an absorbed powerlaw with high energy cutoff plus three Gaussians (same
spectral model as in case of spectral data from the full observation). The
parameters values of second and third Gaussians for second interval were tied
to the values of these components in the first interval. Between the two
intervals, there is a flux change by a factor of two or more. From the
spectral fitting, we found an increase in column density of hydrogen from
$13.3^{+0.7}_{-1.1}\times 10^{-22}$ cm-22 in the first interval to
$22^{+3}_{-2}\times 10^{-22}$ cm-22 in the second interval. The increase in
column density becomes visually apparent when we rescale the spectral model of
the first interval to the flux level of the second interval. In addition, the
equivalent width of 6.4 keV emission line increased from 108 eV in the first
interval to 143 eV in the second interval. Although the spectral index of the
powerlaw component was found match within error bars for both the intervals.
We shall discuss the implication of these spectral changes in the discussion
section.

##  4 Discussion

We report the results from a detailed timing and spectral study of HMXB pulsar
IGR J16320-4751 using all available archival observations taken between 2003
and 2008. We also use a long term Swift-BAT light curve between 2005 and 2020.
Archival observations were used by García et al., (2018) to study the orbital
variation of X-ray emission and column density. They reported high variability
and flaring activity in this source. Also, they derived a correlated evolution
of column density and flux of the Fe $K_{\alpha}$ along the orbit of this
system. Our analysis brought some new results about source behavior.
Specifically, we brought up three interesting observations. These observations
are to do with 1) the timing variability in terms of pulse to pulse variations
in different light curves and average pulse profile variability, 2) the pulse
phase behavior of Fe $K_{\alpha}$ photons, and 3) the presence of soft excess
below 3 keV in the spectra of the source. From the Swift-BAT data of this
source, we found an orbital period of $8.991\pm 0.001$ days which is more
accurate than previously reported values (García et al.,, 2018; Corbet et
al.,, 2005). We discuss the implications of these results in the following
section below.

###  4.1 Timing variability and variable pulse profile

The light curves of IGR J16320-4751 are highly variable. This variability is
in terms of the difference in average count rates from different observations,
flaring activity, and sudden changes in count rate. This seemingly random
variation is characteristic of absorption and accretion from clumpy wind given
by companion (Pradhan et al.,, 2019). In addition to this, pulse to pulse
variations are noticed in 6 out of 11 observations. There are several
instances of pulse to pulse variation for single and multiple pulses together.
During such instances, either pulse is missing or the pulse shape drastically
changed. Spectral comparison of one such instance where good spectral data was
found with an adjoining interval with clear pulsation indicate that these
changes are driven by variable absorption in IGR J16320-4751. The increase in
the absorption column density and equivalent width of the 6.4 keV emission
line indicates that NS is passing through a dense clump of matter when the
pulsation is missing in the light curve. Probably the scattering of soft X-ray
photons by the matter clump causes the smearing of X-ray pulsation. Such
instances are more common in BeXRB sources like GX 304-1, A 0535+26, GRO
J1008-57, and V 0332+53 due to highly variable winds from the companion but
SgXRBs Vela X-1 and OAO 1657-415 has also been found to show such a behavior
(Kretschmar et al.,, 2014; Pradhan et al.,, 2014). Our analysis of IGR
J16320-4751 indicates that such variations can also be observed in highly
obscured SgHMXB sources. The strong timing variability of IGR J16320-4751
leads to variation in the average pulse profile from different observations
even when taken a few days apart. Variability is in pulse shape as well as in
pulse fraction. This makes it difficult to estimate the pulse arrival time
delays using the correlation of pulse profiles. A mass function of $\sim$280
$M_{\odot}$ or orbital period derivative of 5$\times$10-4 ss-1 determined by
using these observations seems unlikely and large phase arrival delays are
driven by intrinsic variability of pulse profiles. If this is indeed true,
hard X-ray pulse profiles and phase lags from NuSTAR/AstroSat observations can
possibly give the orbital solution in this source.

###  4.2 Pulse phase behaviour of the iron line

Fe $K_{\alpha}$ emission lines are ubiquitous features in the spectra of X-ray
binaries with a neutron star. This spectral feature is caused by reprocessing
of hard X-rays from the NS from neutral or partially ionized material in the
system. Candidates for the reprocessing site are the Alfven shell, the
accretion disk (in those sources powered by the Roche lobe overflow),
accretion wakes, the stellar wind, and the surface of the companion star. Our
energy resolved pulse profile revealed a pulsating behavior in 6.2-6.6 keV
energy band. However, the true behavior of Fe $K_{\alpha}$ photons (obtained
by subtracting the proper continuum emission contribution from 6.2-6.6 keV
energy band) is non-pulsating. This is also validated with the pulse phase
resolved spectroscopic analysis. Such a non-pulsating behavior of Fe
$K_{\alpha}$ emission has been reported in some other HMXBs like GX 1+4
(Kotani et al.,, 1999), Vela X-1 (Paul et al., 2002a, ), and 4U 1901+03 (Lei
et al.,, 2009). This result rules out material having an anisotropic
distribution as in the case of Alfven shell, accretion stream, the surface of
the companion star, or a warped structure in the accretion disk. Line energy
of the emission line in 6.2-6.6 keV band obtained from spectral fitting of
different observations lies between 6.38-6.43 keV (see Table 3). These values
correspond to either a low or a moderate ionization state for Fe atoms
(House,, 1969). The highest value of line energy corresponds to Fe XVII.
Hence, the line emitting material is also not located very close to NS. The
distance between the line emitting region and NS can be estimated using the
formula for the ionization parameter, $\zeta$ = L/($N_{H}\times r$), where L
is the luminosity of the source, $N_{H}$ = $\rho\times r$ is column density in
the line of sight towards the source. From the line energy of the emission
line in 6.2-6.6 keV band, we can estimate $log\zeta<2.1$ (Figure 5. of Kallman
et al., (2004)), therefore $\zeta<125$. The unabsorbed bolometric flux in
0.01-100 keV has been estimated to be $4.4\times 10^{-10}$ erg $cm^{-2}$
$s^{-1}$ using a combined spectral fit of the and INTEGRAL (Rodriguez et al.,,
2006). This leads to a luminosity of $6.4\times 10^{35}$ erg $s^{-1}$ for a
distance of 3.5 kpc estimated for this source. Using these values of L and
$\zeta$, we get $r>$1.5-4.3$\times 10^{10}$ cm (corresponding to
$N_{H}$=1.2-3.5$\times 10^{23}$ $cm^{-2}$ obtained from spectral fits) for the
distance between NS and line emitting region. An upper bound on the distance
of the line emitting region comes from the fact that even a symmetrically
distributed material at a large distance can produce pulsed emission due to
difference in light travel time. For $P_{spin}=1305.9$ sec, we get $r<4\times
10^{13}$ cm for line emission to remain unpulsed. The Fe $K_{\alpha}$ line
emitting material is located between these two distance bounds in this source.

###  4.3 Presence of soft excess

Figure 8: Comparison of pulse profile in soft excess band (0.3-3.0 keV) with
the higher energy band (3.0-10.0 keV) for observation #8 and #10 which show
strong soft excess emission.

Soft excess emission is a spectral feature of excess intensities over the
extrapolation of a power law fitted to higher energies. Such emission is seen
in many HMXBs and is explained by 4 possible mechanisms (Hickox et al.,,
2004): (1) emission from the accretion column (2) reprocessing by optically
thick, dense material around the NS (3) thermal emission by collisionally
energized diffuse gas around the NS (4) reprocessing of hard X-rays by diffuse
material around NS. The soft excess emission in the non-flaring duration of
observation #2 was fitted using a blackbody component with a temperature of
0.07 keV (Rodriguez et al.,, 2006). They suggested that the soft excess
emission is due to collisionally energized cloud based on the energy
conservation argument (Rodriguez et al.,, 2006). However, the blackbody
component does not give a good fit for the other observations where such
emission is observed. Another model component of thermal Bremsstrahlung gives
a good fit but the value of plasma temperature seems unusually high ($\sim$200
keV) for such an environment. Other sources like SMC X-1 and LMC X-4 show
plasma temperatures of $\sim$0.1 keV (Paul et al., 2002b, ). A common feature
among all observations of this source showing soft excess is the absence of
pulsation in the 0.3-3.0 keV band (see Figure 8). The absence of pulsation
combined with the fact this source has very low luminosity suggests that the
soft excess is emitted by diffuse gas in this source similar to other sources
like Vela X-1, RX J0101.3-7211, and AX J0101-722 (Hickox et al.,, 2004).
Further, this could be either a collisionally ionized or photoionized gas.
Reprocessing from a diffuse gas cloud has been found to contribute to soft
excess either during the eclipse or just before the eclipse in sources like
Vela X-1 (Schulz et al.,, 2002), 4U 1700-37 (Boroson et al.,, 2003), and 4U
1538-522 (Aftab et al.,, 2019). High-resolution spectral studies in these
sources have shown a large variety of emission features like fluorescent
emission lines from the mid-Z elements and iron, radiative recombination
continuum feature, and lines from H-like, He-like ions. Hence, the soft X-ray
emission of IGR J16320-4751 should be studied with high-resolution spectral
observation with XMM-RGS or Chandra-HETG to confirm the nature of soft excess
and complex composition of the emitting plasma.

Acknowledgements. This work is based on observation data and software tools
from (An ESA Science mission with instruments and contributions directly
funded by ESA member states and USA) and NASA’s High Energy Astrophysics
Science Archive (HEASARC), a service of Goddard space flight center and the
Smithsonian Astrophysical Observatory.

## References

  * Aftab and Paul, (2017) Aftab, N. and Paul, B. (2017).  X-ray reprocessing: Through eclipse spectra of high and low mass X-ray binaries.  In Ness, J.-U. and Migliari, S., editors, The X-ray Universe 2017, page 28.
  * Aftab et al., (2019) Aftab, N., Paul, B., and Kretschmar, P. (2019).  X-Ray Reprocessing: Through the Eclipse Spectra of High-mass X-Ray Binaries with XMM-Newton.  ApJS, 243(2):29.
  * Boldin et al., (2013) Boldin, P. A., Tsygankov, S. S., and Lutovinov, A. A. (2013).  On timing and spectral characteristics of the X-ray pulsar 4U 0115+63: Evolution of the pulsation period and the cyclotron line energy.  Astronomy Letters, 39(6):375–388.
  * Boroson et al., (2003) Boroson, B., Vrtilek, S. D., Kallman, T., and Corcoran, M. (2003).  Chandra Grating Spectroscopy of the X-Ray Binary 4U 1700-37 in a Flaring State.  ApJ, 592(1):516–531.
  * Chaty et al., (2008) Chaty, S., Rahoui, F., Foellmi, C., Tomsick, J. A., Rodriguez, J., and Walter, R. (2008).  Multi-wavelength observations of Galactic hard X-ray sources discovered by INTEGRAL. I. The nature of the companion star.  A&A, 484:783–800.
  * Coleiro et al., (2013) Coleiro, A., Chaty, S., Zurita Heras, J. A., Rahoui, F., and Tomsick, J. A. (2013).  Infrared identification of high-mass X-ray binaries discovered by INTEGRAL.  A&A, 560:A108.
  * Corbet et al., (2005) Corbet, R., Barbier, L., Barthelmy, S., Cummings, J., Fenimore, E., Gehrels, N., Hullinger, D., Krimm, H., Markwardt, C., Palmer, D., Parsons, A., Sakamoto, T., Sato, G., Tueller, J., and Swift-Survey Team (2005).  Swift/BAT Discovery of the Orbital Period of IGR J16320-4751.  The Astronomer’s Telegram, 649.
  * Filippova et al., (2004) Filippova, E. V., Lutovinov, A. A., Shtykovsky, P. E., Revnivtsev, M. G., Burenin, R. A., Arefiev, V. A., Pavlinsky, M. N., and Sunyaev, R. A. (2004).  Broadband Observations of the Transient X-ray Pulsar SAX J2103.5+4545.  Astronomy Letters, 30(12):824–833.
  * Foschini et al., (2004) Foschini, L., Tomsick, J. A., Rodriguez, J., Walter, R., Goldwurm, A., Corbel, S., and Kaaret, P. (2004).  High-energy emission from IGR J16320-4751.  In Schoenfelder, V., Lichti, G., and Winkler, C., editors, 5th INTEGRAL Workshop on the INTEGRAL Universe, volume 552 of ESA Special Publication, pages 247–252.
  * García et al., (2018) García, F., Fogantini, F. A., Chaty, S., and Combi, J. A. (2018).  Spectral evolution of the supergiant HMXB IGR J16320-4751 along its orbit using XMM-Newton.  A&A, 618:A61.
  * Heras et al., (2009) Heras, J. A. Z., Chaty, S., Prat, L., and Rodriguez, J. (2009).  Spectral Evolution Along the Orbit of IGR J16320-4751 with XMM-Newton.  In Rodriguez, J. and Ferrando, P., editors, American Institute of Physics Conference Series, volume 1126 of American Institute of Physics Conference Series, pages 313–315.
  * Hickox et al., (2004) Hickox, R. C., Narayan, R., and Kallman, T. R. (2004).  Origin of the Soft Excess in X-Ray Pulsars.  ApJ, 614:881–896.
  * Ho et al., (2014) Ho, W. C. G., Klus, H., Coe, M. J., and Andersson, N. (2014).  Equilibrium spin pulsars unite neutron star populations.  MNRAS, 437(4):3664–3669.
  * House, (1969) House, L. L. (1969).  Theoretical Wavelengths for Ka-TYPE X-Ray Lines in the Spectra of Ionized Atoms (carbon to Copper).  ApJS, 18:21.
  * Ibarra et al., (2007) Ibarra, A., Matt, G., Guainazzi, M., Kuulkers, E., Jiménez-Bailón, E., Rodriguez, J., Nicastro, F., and Walter, R. (2007).  The XMM-Newton/INTEGRAL monitoring campaign of IGR J16318-4848.  A&A, 465(2):501–507.
  * Jansen et al., (2001) Jansen, F., Lumb, D., Altieri, B., Clavel, J., Ehle, M., Erd, C., Gabriel, C., Guainazzi, M., Gondoin, P., Much, R., Munoz, R., Santos, M., Schartel, N., Texier, D., and Vacanti, G. (2001).  XMM-Newton observatory. I. The spacecraft and operations.  A&A, 365:L1–L6.
  * Kallman et al., (2004) Kallman, T. R., Palmeri, P., Bautista, M. A., Mendoza, C., and Krolik, J. H. (2004).  Photoionization Modeling and the K Lines of Iron.  ApJS, 155:675–701.
  * Kotani et al., (1999) Kotani, T., Dotani, T., Nagase, F., Greenhill, J. G., Pravdo, S. H., and Angelini, L. (1999).  ASCA and Ginga Observations of GX 1+4.  ApJ, 510(1):369–378.
  * Kretschmar et al., (2014) Kretschmar, P., Marcu, D., Kühnel, M., Klochkov, D., Pottschmidt, K., Staubert, R., Wilson-Hodge, C. A., Jenke, P. A., Caballero, I., and Fürst, F. (2014).  Pulse-to-pulse variations in accreting X-ray pulsars.  In European Physical Journal Web of Conferences, volume 64 of European Physical Journal Web of Conferences, page 06012.
  * Lebrun et al., (2001) Lebrun, F., Leray, J.-P., Limousin, O., Blondel, C., Laurent, P., Goldoni, P., Goldwurm, A., and Sauvageon, A. (2001).  ISGRI: The CdTe gamma-ray camera onboard INTEGRAL.  In Gimenez, A., Reglero, V., and Winkler, C., editors, Exploring the Gamma-Ray Universe, volume 459 of ESA Special Publication, pages 591–594.
  * Lei et al., (2009) Lei, Y.-J., Chen, W., Qu, J.-L., Song, L.-M., Zhang, S., Lu, Y., Zhang, H.-T., and Li, T.-P. (2009).  Phase-Resolved Spectral Analysis of 4U 1901+03 During Its Outburst.  ApJ, 707(2):1016–1022.
  * Liu et al., (2000) Liu, Q. Z., van Paradijs, J., and van den Heuvel, E. P. J. (2000).  A catalogue of high-mass X-ray binaries.  A&AS, 147:25–49.
  * Lutovinov et al., (2005) Lutovinov, A., Rodriguez, J., Revnivtsev, M., and Shtykovskiy, P. (2005).  Discovery of X-ray pulsations from IGR J16320-4751 = AX J1631.9-4752.  A&A, 433(2):L41–L44.
  * (24) Paul, B., Dewangan, G. C., Sako, M., Kahn, S. M., Paerels, F., Liedahl, D., Wojdowski, P., and Nagase, F. (2002a).  Pulse Phase Resolved Spectroscopy of the Fe Kα Line of Vela X-1 with Chandra HETG.  In Ikeuchi, S., Hearnshaw, J., and Hanawa, T., editors, 8th Asian-Pacific Regional Meeting, Volume II, pages 355–356.
  * (25) Paul, B., Nagase, F., Endo, T., Dotani, T., Yokogawa, J., and Nishiuchi, M. (2002b).  Nature of the Soft Spectral Component in the X-Ray Pulsars SMC X-1 and LMC X-4.  ApJ, 579(1):411–421.
  * Pradhan et al., (2018) Pradhan, P., Bozzo, E., and Paul, B. (2018).  Supergiant fast X-ray transients versus classical supergiant high mass X-ray binaries: Does the difference lie in the companion wind?  A&A, 610:A50.
  * Pradhan et al., (2019) Pradhan, P., Bozzo, E., Paul, B., Manousakis, A., and Ferrigno, C. (2019).  Probing Clumpy Wind Accretion in IGR J18027-2016 with XMM-Newton.  ApJ, 883(2):116.
  * Pradhan et al., (2014) Pradhan, P., Maitra, C., Paul, B., Islam, N., and Paul, B. C. (2014).  Variations in the pulsation and spectral characteristics of OAO 1657-415.  MNRAS, 442:2691–2700.
  * Rodriguez et al., (2006) Rodriguez, J., Bodaghee, A., Kaaret, P., Tomsick, J. A., Kuulkers, E., Malaguti, G., Petrucci, P.-O., Cabanac, C., Chernyakova, M., Corbel, S., Deluit, S., Di Cocco, G., Ebisawa, K., Goldwurm, A., Henri, G., Lebrun, F., Paizis, A., Walter, R., and Foschini, L. (2006).  INTEGRAL and XMM-Newton observations of the X-ray pulsar IGR J16320-4751/AX J1631.9-4752.  MNRAS, 366:274–282.
  * (30) Rodriguez, J., Del Santo, M., Lebrun, F., Belanger, G., Capitanio, M., Cadolle-Bel F., David, P., Foschini, L., Goldoni, P., Goldwurm, A., Gros, A., Laurent, P., Paizis, A., Paul, J., Terrier, R., Shaw, S. E., and Ubertini, P. (2003a).  First results from the IBIS/ISGRI data obtained during the Galactic Plane Scan. II. The Vela region.  A&A, 411:L373–L376.
  * (31) Rodriguez, J., Tomsick, J. A., Foschini, L., Walter, R., Goldwurm, A., Corbel, S., and Kaaret, P. (2003b).  An XMM-Newton observation of IGR J16320-4751 = AX J1631.9-4752.  A&A, 407:L41–L45.
  * Schulz et al., (2002) Schulz, N. S., Canizares, C. R., Lee, J. C., and Sako, M. (2002).  The Ionized Stellar Wind in Vela X-1 during Eclipse.  ApJ, 564(1):L21–L25.
  * Sugizaki et al., (2001) Sugizaki, M., Mitsuda, K., Kaneda, H., Matsuzaki, K., Yamauchi, S., and Koyama, K. (2001).  Faint X-Ray Sources Resolved in the ASCA Galactic Plane Survey and Their Contribution to the Galactic Ridge X-Ray Emission.  ApJS, 134:77–102.
  * Titarchuk, (1994) Titarchuk, L. (1994).  Generalized Comptonization Models and Application to the Recent High-Energy Observations.  ApJ, 434:570.
  * Tomsick et al., (2003) Tomsick, J. A., Lingenfelter, R., Walter, R., Rodriguez, J., Goldwurm, A., Corbel, S., and Kaaret, P. (2003).  Igr J16320-4751.  IAU Circ., 8076.
  * Torrejón et al., (2010) Torrejón, J. M., Schulz, N. S., Nowak, M. A., and Kallman, T. R. (2010).  A Chandra Survey of Fluorescence Fe Lines in X-ray Binaries at High Resolution.  ApJ, 715:947–958.
  * Walter et al., (2015) Walter, R., Lutovinov, A. A., Bozzo, E., and Tsygankov, S. S. (2015).  High-mass X-ray binaries in the Milky Way. A closer look with INTEGRAL.  A&A Rev., 23:2.
  * Zurita et al., (2006) Zurita, C., Torres, M. A. P., Steeghs, D., Rodríguez-Gil, P., Muñoz-Darias, T., Casares, J., Shahbaz, T., Martínez-Pais, I. G., Zhao, P., Garcia, M. R., Piccioni, A., Bartolini, C., Guarnieri, A., Bloom, J. S., Blake, C. H., Falco, E. E., Szentgyorgyi, A., and Skrutskie, M. (2006).  The 2005 Outburst of the Halo Black Hole X-Ray Transient XTE J1118+480.  ApJ, 644(1):432–438.


= 0.16em

# How should a small country respond to climate change?

A. D. J. Haymet  $|$ ${}^{\textsf{\footnotesize{\mbox{?},\penalty
1000\mbox{?}}}}$

(Received April 12, 2023, in final form April 23, 2023)

###### Abstract

Responses to the global climate crisis often focus on the largest current
emitters of greenhouse gases. However, analysis shows that about a third of
emissions come from a collection of small emitters, each contributing one- to
two-percent of the total additional CO2 injected into the communal atmosphere.
Attempts to hold global warming to less than 1.5℃ cannot succeed without also
reducing emissions from these small countries.

Key words: climate change, greenhouse gases, small country response

###### Abstract

Ãîâîðÿè ïðî ãëîáàëüíó êëìàòèíó êðèçó, àñòî çîñåðåäæóþòüñÿ íà êðà¿íàõ, ÿê íà
òåïåðøíé àñ âèêèäàþòü íàéáëüøó êëüêñòü ïàðíèêîâèõ ãàçâ. Îäíàê àíàëç ïîêàçó, ùî
áëèçüêî òðåòèíè âèêèäâ çóìîâëþþòü ìàë êðà¿íè, íà äîëþ êîæíî¿ ç ÿêèõ ïðèïàäà âä
îäíîãî äî äâîõ âäñîòêâ çàãàëüíîãî íàäëèøêîâîãî CO2, ùî âèêèäàòüñÿ â àòìîñôåðó.
Ñïðîáè âòðèìàòè ãëîáàëüíå ïîòåïëííÿ â ìåæàõ 1.5℃ íå ìîæóòü áóòè óñïøíèìè áåç
çìåíøåííÿ âèêèäâ öèõ ìàëèõ êðà¿í.

Ключов слова: çìíè êëìàòó, ïàðíèêîâ ãàçè, ðåàêöÿ ìàëî¿ êðà¿íè

What should small countries deeply affected by GHG (greenhouse gas) emissions
do? How do they convince large countries to drastically reduce emissions,
especially if they themselves mine the products which — when used — greatly
add to global warming? For example, to save the Great Barrier Reef, a major
tourist attraction and biodiverse ecosystem, Australia must convince the
northern hemisphere to stop emitting CO2 and other heat-trapping greenhouse
gases as quickly as possible. Since 1955, by actual measurements up to
December 2022 (not “calculations” or “theory”) humans have dumped over 345
($\pm 2$) zettajoules of extra heat into the earth system [1], almost all of
it into the ocean. That extra heat fuels the weather every single day. It
perverts our winds and rainfall. It supercharges cyclones. It bleaches coral
reefs. It causes sea-level rise. It disturbs our fisheries and agriculture.
And it creates a higher probability of extreme bushfire events.

There is no doubt that the damage comes from this excess heat. So where does
the heat come from? That is also no mystery. Fortunately, since 1965 (at
least) we have known exactly. We humans have increased the atmospheric CO2
concentration by 50%, to over 410 ppm. That CO2 plus methane and other
greenhouse gases (GHG) (assisted by some black soot) trap more heat than
nature intended in the atmosphere, and that heat is rapidly transferred to the
ocean, due to its massively higher heat capacity. Since 1958, no other
explanation has survived scientific scrutiny.

About 92.5% of the heat stays in the ocean [2]. Only about 2% stays in the
atmosphere, but since that is where we live, we tend to focus on “global
warming” of the atmosphere. The massive heating of the ocean is much more
important and much more damaging. Additional small percentages of that heat
melt glaciers formerly grounded on land, and also heat up the soil and land
mass directly.

In many ways it is a great relief that we know precisely what is causing this
extra heat. Because knowledge is power. We can reverse it, albeit agonizingly
slowly. Imagine if our scientific knowledge of ocean warming was not so
precise! We would be suffering in terrible uncertainty about how to respond to
the current crisis.

But we know exactly what to do. To save Australia’s fisheries and farms, its
reefs and its forests and its very lifestyle, we need to stop the production
of this extra heat, and we can do that by steadily quickly reducing the global
output of greenhouse gases, and eventually become “carbon neutral”.

Thanks to scientist Susan Solomon [3], we have known for a decade how long
this reversal takes. It is agonizingly long. That which we sullied in 50–100
years takes a thousand years to undo. But that is NOT an argument to do
nothing, but rather an exhortation — a commandment if you will — to start
reductions today. Many jurisdictions, including Sweden and California, and
major companies all aim to reach carbon neutrality by 2030. Australia is
blessed with copious renewable resources, more than enough to make hydrogen
from renewables, and superb engineering & scientific talent. We are indeed the
“lucky country” in responding to this excess heat challenge.

Imagine if we had acted upon the rigorous scientific reports of 1965, the
Lyndon Johnson “Presidents Science Advisory Committee Report on Atmospheric
Carbon Dioxide” [4], and the 1979 report chaired by Charney [5], and had
steadily reduced our GHG emissions starting 55 years ago. Our current observed
global warming of 1.1℃ would be vastly reduced, and our time to reverse the
harm vastly shorter. The currently observed bias towards formerly unlikely
disastrous events would be much reduced. Alas we cannot go back and implement
their wisdom over the last 55 years. But we can start today. In another 55
years, the next two generations will at least be thanking us for starting
today, as much as they will curse us for not acting in 1965.

Why is the job of Australia to convince the northern hemisphere to reduce
their GHG output? Because their pollution eventually reaches us. The data in
figure 2 below, documents the difference in CO2 concentration between northern
and southern hemispheres, since 1958.

Figure 1: (Colour online) Pollution from the northern hemisphere. Black dots:
difference of the monthly average atmospheric carbon dioxide concentration at
Mauna Loa Observatory, Hawaii and South Pole, Antarctica versus time where CO2
concentration is in parts per million (ppm). Black curve: spline fit to the
CO2 data. Data from Scripps CO2 Program. URL
https://scrippsco2.ucsd.edu/graphics_gallery/mauna_loa_and_south_pole/mauna_loa_and_south_pole_difference.html

Figure 2: (Colour online) Black curve: Monthly average atmospheric carbon
dioxide concentration versus time at Mauna Loa Observatory, Hawaii (20∘N,
156∘W) where CO2 concentration is in parts per million in the mole fraction
(ppm). Monthly data are shown as dots and connected with straight lines. The
record is distinguished by its pronounced seasonal cycle. Red Curve: Monthly
average atmospheric carbon dioxide concentration versus time at the South
Pole, Antarctica, where CO2 concentration is in parts per million (ppm).
Monthly data points are shown as dots and connected approximately by a smooth
curve. Data from Scripps CO2 Program. URL
https://scrippsco2.ucsd.edu/graphics_gallery/mauna_loa_and_south_pole/mauna_loa_and_south_pole.html

Figure 3: (Colour online) Annual CO2 emissions from fuel combustion by country
or country group. From IEA data.

Figure 1 arises from the famous Keeling CO2 measurements at Mauna Loa
Observatory, Hawaii, combined with the South Pole observatory in Antarctica,
shown in figure 2. In figure 2, both data sets increase over the decades, due
to human emissions, but the northern hemisphere is increasing faster. It takes
quite a few years now for the extra CO2 in the northern hemisphere to make it
across the natural atmospheric barriers at the equator to the southern
hemisphere. How should Australia go about this vital task?

Australia has a choice of two basic tactics to get the northern hemisphere to
stop emitting:

  * •

“Do as I do”,

  * •

“Do as I say”.

“Do as I do” means setting aggressive reduction targets for Australia,
including the consequential emissions from its exports. It means avoiding
accounting tricks from a long-forgotten late-night bargaining session in Kyoto
by former Australian Environment Minister Robert Hills righthand person.

“Do as I say” means… well, what does it mean? How does one go about convincing
the many dozens of global emitting countries to do what Australians themselves
are unwilling to do? As shown in figure 2, far from reducing global emissions,
last year GHG emissions increased and a record rate. We are making a
horrendous problem even worse, each day, every day, even though we know
exactly how to fix it. We can make electricity without generating CO2 or
releasing methane. The same for transportation. And farming practices. We
still have work to do in metallurgy and cement manufacture and in other
industrial processes, but we have all the ingenuity and motivation we need.

Just how many countries does Australia have to convince? A lot. More than one
third of total global emissions come from countries who are “one-to-two
percent” contributors, just like Australia. Roughly 40 countries emit between
2% and 0.1% of the global total. That third cannot be ignored, even if strong
action were to be taken by the current six heaviest emitters China, USA,
India, Russia, Japan and Germany (60% of total emissions). Consideration of
emissions per capita, and accumulated historical emissions per capita, rather
than just current total per country, lead to an even greater increased
responsibility for Australia.

## Conclusions

Over the last 15 years there have been some great Foreign Ministers of
Australia. Sometimes they have been sent out to do their job with proven
mechanisms in place to reduce Australia’s CO2 pollution, like the emissions
reduction scheme former PM Julia Gillard promised in the 2010 election, and
legislated successfully. At other times Australia has sent Foreign Ministers
out into the world with their arms tied behind their back. Australia must have
a consistent and effective Australian policy, and a good time to start is —
now.
Tony Haymet PhD FTSE is a scientist, business founder and owner, and Chair of
the Antarctic Science Foundation. He has served as Chief of CSIRO Marine &
Atmospheric Research, and Director of Scripps Institution of Oceanography.

## References

  * [1] URL https://climate.nasa.gov/vital-signs/ocean-warming/.
  * [2] URL https://www.ipcc.ch/report/ar5/wg1/.
  * [3] Solomon S., Daniel J. S., Sanford T. J., Friedlingstein P., Proc. Natl. Acad. Sci. U.S.A., 2010, 107, No. 43, 18354–18359, doi:10.1073/pnas.1006282107.
  * [4] 1965 Presidents Science Advisory Committee Report on Atmospheric Carbon Dioxide (USA), URL https://www.climatefiles.com/climate-change-evidence/presidents-report-atmospher-carbon-dioxide/.
  * [5] National Research Council, 1979, Carbon Dioxide and Climate: A Scientific Assessment, The National Academies Press, Washington, doi:10.17226/12181.

Ukrainian 0

ßê ìàëé êðà¿í ñëä ðåàãóâàòè íà çìíè êëìàòó? À. Ä. Ä.
Õàéìåò${}^{\textsf{\footnotesize{\mbox{?},\penalty 1000\mbox{?}}}}$

${}^{\textsf{\footnotesize 1}}$ Àíòàðêòèíà íàóêîâà ôóíäàöÿ, 203 åííåë Õàéâåé,
Êíãñòîí, Òàñìàíÿ, Àâñòðàëÿ ${}^{\textsf{\footnotesize 2}}$ Îêåàíîãðàôíèé
íñòèòóò Ñêðèïïñà, Êàëôîðíéñüêèé óíâåðñèòåò â Ñàí-Äãî, Ñàí Äãî, Êàëôîðíÿ, ÑØÀ