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. The upper panel shows the surface density of stars Figurein the X-Z plane while the lower panel shows the stars which have been identified as groups. | The upper panel shows the surface density of stars in the X-Z plane while the lower panel shows the stars which have been identified as groups. |
It can be seen that our scheme is very successful in detecting the structures. | It can be seen that our scheme is very successful in detecting the structures. |
Even faint structures that are not easily visible in the surface density maps are detected by the group finder. | Even faint structures that are not easily visible in the surface density maps are detected by the group finder. |
EnLink is a hierarchical group finder and the reported groups obey a parent child relationship, with one master group being at the top. | EnLink is a hierarchical group finder and the reported groups obey a parent child relationship, with one master group being at the top. |
In our case this master group represents the smooth component of the halo and the rest of the groups are considered as substructures within it. | In our case this master group represents the smooth component of the halo and the rest of the groups are considered as substructures lying within it. |
Note, that since we analyze magnitude limited lyingsurveys we expect two density peak even in a smooth halo: one at the galactic center and the other at the location of the sun due to the presence of a large number of low stars. | Note, that since we analyze magnitude limited surveys we expect two density peak even in a smooth halo: one at the galactic center and the other at the location of the sun due to the presence of a large number of low luminosity stars. |
Hence we ignore all groups whose density peaks lie luminositywithin 5 kpc of sun or the galactic center. | Hence we ignore all groups whose density peaks lie within 5 kpc of sun or the galactic center. |
When we apply EnLink to our synthetic data sets, because we know the progenitor satellites from which the stars originally came from, we can make some quantitative assessment of how well the group finder has performed. | When we apply EnLink to our synthetic data sets, because we know the progenitor satellites from which the stars originally came from, we can make some quantitative assessment of how well the group finder has performed. |
Following Paper I, we define purity as the maximum fraction | Following Paper I, we define as the maximum fraction |
whereas the effective radius itself can be derived from empirical fits to (Mamon Lookas 2005a) as Lio). where £40=hzp/(10 L3. Lg2Mosf(bgYp). and Lg and Mo; are in units of L. and M. respectively. | whereas the effective radius itself can be derived from empirical fits to (Mamon okas 2005a) as $\log(h_{70}r_e) = 0.34 + 0.54 \log L_{10} + 0.25 (\log L_{10})^2$ , where $L_{10} = h_{70}^2 L_B/(10^{10} L_{\odot})$ , $L_B=M_{vir}/(b_{\Upsilon}\overline{\Upsilon}_B)$, and $L_B$ and $M_{vir}$ are in units of $_{\odot}$ and $_{\odot}$, respectively. |
Finally. we can calculate the central gas density po by imposing where we have assumed that all matter is comprised of DM. gas; and stars only (ic. we have neglected other forms of matter. such as dust. and the presence of a supermassive central black hole). and that p(r) is given by equation or when stars are included in the gravitational potential. and by equation or when stars are excluded. | Finally, we can calculate the central gas density $\rho_0$ by imposing where we have assumed that all matter is comprised of DM, gas, and stars only (i.e. we have neglected other forms of matter, such as dust, and the presence of a supermassive central black hole), and that $\rho(r)$ is given by equation or when stars are included in the gravitational potential, and by equation or when stars are excluded. |
Note. however. that the eas virial mass is the same with and without the stellar component in the gravitational potential. since stars are included in the total mass budget in both cases. | Note, however, that the gas virial mass is the same with and without the stellar component in the gravitational potential, since stars are included in the total mass budget in both cases. |
We note that à very small number of parameters are needed to compute the equilibrium profiles for the gas. | We note that a very small number of parameters are needed to compute the equilibrium profiles for the gas. |
These are. for the DM halo system. the redshift z. the (total) virial mass οι of the system. the fitting parameters ey and a of the empirical relation between concentration and total virial mass. and the barvon fraction 6, relative to the universal value. | These are, for the DM halo system, the redshift $z$ , the (total) virial mass $M_{vir}$ of the system, the fitting parameters $c_0$ and $\alpha$ of the empirical relation between concentration and total virial mass, and the baryon fraction $b_b$ relative to the universal value. |
For the stars. the parameters are the galactie B-bancl mass-to-light ratio Ys» (or. alternatively. the universal B-band mass-to-light ratio Yg and the D-band mass-to-light ratio bias by). the stellar B-band mass-to-light ratio Tyo. and the Dehnen parameter v. | For the stars, the parameters are the galactic B-band mass-to-light ratio $\Upsilon_B$ (or, alternatively, the universal B-band mass-to-light ratio $\overline{\Upsilon}_B$ and the B-band mass-to-light ratio bias $b_{\Upsilon}$ ), the stellar B-band mass-to-light ratio $\Upsilon_{*,B}$, and the Dehnen parameter $\psi$. |
Finally. for the gas. the parameters are the mean molecular weight so the polvtropic index E. and the central temperature Tu (or. alternatively. the ratio between central density and. pressure po £1). | Finally, for the gas, the parameters are the mean molecular weight $\mu$, the polytropic index $\Gamma$, and the central temperature $T_0$ (or, alternatively, the ratio between central density and pressure $\rho_0/P_0$ ). |
In our fiducial model. we use the Maimon Lookas (2005b) mass-to-light ratio values and choose by=0.25641. Tex=390 AL. fL. (so that Yg=100 M. 7L. precisely). and Ἐνο26.5 AL. fL.. | In our fiducial model, we use the Mamon okas (2005b) mass-to-light ratio values and choose $b_{\Upsilon}=0.25641$, $\overline{\Upsilon}_B=390$ $_{\odot}/$ $_{\odot}$ (so that $\Upsilon_B=100$ $_{\odot}/$ $_{\odot}$ precisely), and $\Upsilon_{*,B}=6.5$ $_{\odot}/$ $_{\odot}$. |
As for the halo parameters. we choose a svstem with universal barvon fraction (b,= 1) at redshift z=0.5. a (total) virial mass Mo;=107 M... and the e.;,-Al.;, relation parameters ey=9 and à=0.19 (Bullock et al. | As for the halo parameters, we choose a system with universal baryon fraction $b_b=1$ ) at redshift $z=0.5$, a (total) virial mass $M_{vir}=10^{12}$ $_{\odot}$, and the $c_{vir}$ $M_{vir}$ relation parameters $c_0=9$ and $\alpha=-0.13$ (Bullock et al. |
2001). | 2001). |
Finally. we fix the central temperature of the gas to be Apdo=0.6 keV (Diehl Statler 2008). the mean molecular weight of the gas to be ji=0.62. the polvtropic index to be P=5/3. ane the Dehnen parameter to be c=1 (Llernquist 1990). | Finally, we fix the central temperature of the gas to be $k_BT_0=0.6$ keV (Diehl Statler 2008), the mean molecular weight of the gas to be $\mu=0.62$, the polytropic index to be $\Gamma=5/3$, and the Dehnen parameter to be $\psi=1$ (Hernquist 1990). |
IE these quantities are listed in "Fable 1.. in addition to the parameters calculated throughout this section. | All these quantities are listed in Table \ref{tab:Parameters}, in addition to the parameters calculated throughout this section. |
The choice of the central temperature Zo is particularly important. since it greatly. inlluences the values of Aypy ane Ap. | The choice of the central temperature $T_0$ is particularly important, since it greatly influences the values of $\Delta_{NFW}$ and $\Delta_D$. |
M we use the (cluster and. group) mass-tempoerature (ALT) relation (e.g. Sanderson et al. | If we use the (cluster and group) mass-temperature (M-T) relation (e.g. Sanderson et al. |
2003) or. rather. extrapolate the ALP relation down to the mass ranges tvpical of early-type galaxies. we obtain e.g. Apdoz0.2 keV for a system with AG,=107 M. | 2003) or, rather, extrapolate the M-T relation down to the mass ranges typical of early-type galaxies, we obtain e.g. $k_BT_0 \simeq 0.2$ keV for a system with $M_{vir}=10^{12}$ $_{\odot}$. |
However. the extrapolation of the M-T. relation is not without risk. as there is some evidence for a gradua steepening of the relation. with decreasing mass (Sanderson ct al. | However, the extrapolation of the M-T relation is not without risk, as there is some evidence for a gradual steepening of the relation, with decreasing mass (Sanderson et al. |
2003). | 2003). |
Using individual observations of isolated. elliptica galaxies. the central temperatures measured. are usually slightly higher. ranging from Apdo0.5 keV (e.g. Memola ct al. | Using individual observations of isolated elliptical galaxies, the central temperatures measured are usually slightly higher, ranging from $k_BT_0 \sim 0.5$ keV (e.g. Memola et al. |
2009) to ketoI keV (e.g. O'Sullivan Ponman 2004: O'Sullivan. Sanderson Ponman 2007).Finally. DiehlStatler (2008) give a convincing argument for acentral temperature &gZu0.6 keV. using observations of 36 normal earlv-tyvpe galaxies. | 2009) to $k_BT_0 \sim 1$ keV (e.g. O'Sullivan Ponman 2004; O'Sullivan, Sanderson Ponman 2007).Finally, DiehlStatler (2008) give a convincing argument for acentral temperature $k_BT_0 \sim 0.6$ keV, using observations of 36 normal early-type galaxies. |
For this reason. here we study cases with three dilferent central temperatures: Ago=0.2.0.6.1 keV. In Fig. | For this reason, here we study cases with three different central temperatures: $k_BT_0=0.2,0.6,1$ keV. In Fig. |
1. we plot the density (top panels) and enclosed mass (middle panels) profiles of DM. stars. and isothermal gas. for | \ref{fig:Dependence_on_T} we plot the density (top panels) and enclosed mass (middle panels) profiles of DM, stars, and isothermal gas, for |
the main luminosity criteria are the strengths of the lines. | the main luminosity criteria are the strengths of the lines. |
Strong AA 4089 4116 lines are clearly. visible on the wings of H9. | Strong $\lambda\lambda$ 4089 4116 lines are clearly visible on the wings of $\delta$. |
The ratio A4089 ~ A4121 is indicative of a giant star (Walborn Fitzpatrick 1990). | The ratio $\lambda$ 4089 $\sim$ $\lambda$ 4121 is indicative of a giant star (Walborn Fitzpatrick 1990). |
Phe strength of A4009 aargues against a spectral tvpe earlier than O9.5. while the absence of and the presence of make the object earlier than BO.5. | The strength of $\lambda$ 4009 argues against a spectral type earlier than O9.5, while the absence of and the presence of make the object earlier than B0.5. |
Since both A 4541 aanel the triplet can be seen in lower-resolution spectra (e.g... Clark et al. | Since both $\lambda$ 4541 and the triplet can be seen in lower-resolution spectra (e.g., Clark et al. |
1998) we believe that BOlLLe is the most appropriate classification. though the previously accepted O9.71LHIe cannot be ciscarclect. | 1998) we believe that B0IIIe is the most appropriate classification, though the previously accepted O9.7IIIe cannot be discarded. |
Llowarth (1983) made an analysis of eight short wavelength and five long wavelength LUE spectra of LSL 61 303. | Howarth (1983) made an analysis of eight short wavelength and five long wavelength IUE spectra of LSI $^\circ$ 303. |
By Ilattening the extinction bump he derives (οὐ1)=0.75x 0.1. | By flattening the extinction bump he derives $E(B-V)=0.75\pm0.1$ . |
Such a large reddening is inconsistent. with a distance of only ~200 pe (Ishida 1969). | Such a large reddening is inconsistent with a distance of only $\sim 200$ pc (Ishida 1969). |
Lt is therefore important to determine the true value of the reddening to the svstem. | It is therefore important to determine the true value of the reddening to the system. |
We do this using two methods — the sodium D» line in our spectra and the strength of optical cilluse interstellar bancs. | We do this using two methods – the sodium $_2$ line in our spectra and the strength of optical diffuse interstellar bands. |
The first method. we employ. uses the correlation between the strength of the interstellar Sodium D» line and extinction derived. by Hobbs. (1974). | The first method we employ uses the correlation between the strength of the interstellar Sodium $_2$ line and extinction derived by Hobbs (1974). |
. We have. rederived £(BV) values for all of Hobbs (1974) objects for which he quotes Na De equivalent widths using spectral classifications and DVY colours from the Bright Star Catalogue (Llollleit Jaschek 1982) and the intrinsic colours of Popper (1980) [or dwarls ancl giants and Johnson (1966) for supergiants. | We have rederived $E(B-V)$ values for all of Hobbs (1974) objects for which he quotes Na $_2$ equivalent widths using spectral classifications and $B-V$ colours from the Bright Star Catalogue (Hoffleit Jaschek 1982) and the intrinsic colours of Popper (1980) for dwarfs and giants and Johnson (1966) for supergiants. |
The correlation in plotted. in Fig.4. | The correlation in plotted in Fig.4. |
The scatter in this diagram is large. and shows that this technique will not be particularly accurate in deriving £(2V). | The scatter in this diagram is large, and shows that this technique will not be particularly accurate in deriving $E(B-V)$. |
Phe dataset does not extend bevond £(BV)~0.7. and is sparse bevond LBV)—0.3. | The dataset does not extend beyond $E(B-V)\sim0.7$, and is sparse beyond $E(B-V)\sim0.3$. |
Recognising the limitations of Fig. | Recognising the limitations of Fig. |
4 however. we may still attempt to use the Na D» line to determine. the extinetion to the svstem. | 4 however, we may still attempt to use the Na $_2$ line to determine the extinction to the system. |
Pwo spectra of LSL 61° 303 were obtained in the region covering theNa D lines on the nights of 1994 March. 26 ancl 27 using the RBS of the JINT. | Two spectra of LSI $^\circ$ 303 were obtained in the region covering theNa D lines on the nights of 1994 March 26 and 27 using the RBS of the JKT. |
The mean Na Do EW measured was 650490 mA... | The mean Na $_2$ EW measured was $650\pm90$ . |
From Figure 4 | From Figure 4 |
lline. | line. |
The ratio of the recombination coefficients is not very sensitive to the assumedΤο. | The ratio of the recombination coefficients is not very sensitive to the assumed. |
In the case of Barnard's Loop we are observing a relatively thin shell seen nearly edge-on, so that the surface brightness will be enhanced by a factor of g above that for a slab viewed face on. | In the case of Barnard's Loop we are observing a relatively thin shell seen nearly edge-on, so that the surface brightness will be enhanced by a factor of g above that for a slab viewed face on. |
We call the quantity the limb-brightening correction. | We call the quantity the limb-brightening correction. |
In this thin shell model the surface brightness distribution would start at the maximum radius μας, rapidly rise to a peak value at angle corresponding to the inner radius Ry, then slowly decrease to the value that would apply along the line of sight through the ionizing star. | In this thin shell model the surface brightness distribution would start at the maximum radius $\rm _{max}$, rapidly rise to a peak value at angle corresponding to the inner radius $\rm _{peak}$, then slowly decrease to the value that would apply along the line of sight through the ionizing star. |
The enhancement of the surface brightness will be g=2( R2,,-R2,,)?/(Ra - Rpeak). | The enhancement of the surface brightness will be g=2( $\rm _{max}^{2} $ $\rm _{peak}^{2}$ $^{1/2}$ $\rm _{max}$ - $\rm _{peak}$ ). |
Using the dimensions for the line from tthrough our CTIO-2008 observations (R;,44— 7:007 and Rpeak=6°001), the geometric correction factor is g—7.03, that is, the observed peak surface brightness will be enhanced by a factor of 7.03. | Using the dimensions for the line from through our CTIO-2008 observations $\rm _{max}$ 07 and $\rm _{peak}$ 01), the geometric correction factor is g=7.03, that is, the observed peak surface brightness will be enhanced by a factor of 7.03. |
One can derive the flux of ionizing photons (¢(H)) from S(H8)) after the former has been corrected for limb brightening. | One can derive the flux of ionizing photons $\phi$ (H)) from ) after the former has been corrected for limb brightening. |
Using our observed surface brightness (5.3x10° steradian!)), the geometric correction factor of 7.03 yields a flux of ionizing photons ¢(H)=8.1 x 10" photons cm"? s~!. | Using our observed surface brightness $5.3\times 10^{6}$ ), the geometric correction factor of 7.03 yields a flux of ionizing photons $\phi$ (H)=8.1 x $^{7}$ photons $^{-2}$ $^{-1}$. |
It is hard to ascribe a probable error to this number. | It is hard to ascribe a probable error to this number. |
The observational uncertainty is less than10%. | The observational uncertainty is less than. |
. The uncertainties in the derived results that are due to applying a simple thin-shell model to Barnard's Loop, which has some internal structure, are probably significantly larger. | The uncertainties in the derived results that are due to applying a simple thin-shell model to Barnard's Loop, which has some internal structure, are probably significantly larger. |
As we will see in the next section, this flux and its wavelength distribution determine the expected photoionization structure and observed line ratios within Barnard's Loop. | As we will see in the next section, this flux and its wavelength distribution determine the expected photoionization structure and observed line ratios within Barnard's Loop. |
Having calculated the ionizing flux in a given sample, one can calculate the total luminosity in LyC photons (Q(H) with the units photons s!) for the ionizing star(s) by the relationQ(H)=4m7 R?9(H). | Having calculated the ionizing flux in a given sample, one can calculate the total luminosity in LyC photons (Q(H) with the units photons $^{-1}$ ) for the ionizing star(s) by the relation$\pi$ $^{2}$$\phi$ (H). |
Adopting the average of the maximum and peak surface brightness radii of 1.56x107° cm yields Q(H)=2.5x1045 photons s. | Adopting the average of the maximum and peak surface brightness radii of $1.56\times 10^{20}$ cm yields $2.5\times 10^{49}$ photons $^{-1}$ . |
wavelengths bands. | wavelengths bands. |
Almost all the data presented here have already been published previously. | Almost all the data presented here have already been published previously. |
Similarly given a model of primary non-Gaussianity one can construct a theoretical model for computation of DisΕΠΗ (see Munshi Heavens 2009 for more about various models and construction of optimal estimators). | Similarly given a model of primary non-Gaussianity one can construct a theoretical model for computation of $B_{ll_1l_2}^{prim}$ (see Munshi Heavens 2009 for more about various models and construction of optimal estimators). |
While study of primary non-Gaussianity is important in its own right for the study of secondaries they can confuse the study. | While study of primary non-Gaussianity is important in its own right for the study of secondaries they can confuse the study. |
Similar results hold at the level of the skew spectrum. | Similar results hold at the level of the skew spectrum. |
A more general treatment based on Fisher analysis of multiple bispectra is presented in the subsequent sections. | A more general treatment based on Fisher analysis of multiple bispectra is presented in the subsequent sections. |
In addition to various sources Mentioned above. second-order corrections to the gravitational potential through gravitational instability too can also act as a source of secondary non-Gaussianity (Munshi.Souradeep&Starobinsky1995). | In addition to various sources mentioned above, second-order corrections to the gravitational potential through gravitational instability too can also act as a source of secondary non-Gaussianity \citep{MuSoSt95}. |
. Starting from Babich(20€ a complete MN of bispectrum in M ee of partial sky coverage and inhomogeneous noise was developedby various authors (Babich2005:"Creminellietal.2006:Yadavum). | Starting from \citet{Babich} a complete analysis of bispectrum in the presence of partial sky coverage and inhomogeneous noise was developedby various authors \citep{Babich,Crem06,Yadav08}. |
A Specific form for a bispectrum estimator was introduced which is both unbiased and optimal. | A specific form for a bispectrum estimator was introduced which is both unbiased and optimal. |
This was further -- and used by Sun.&Dore(2007) for lensing reconstruction and by for general bispectrum analysis. | This was further developed and used by \citet{SmZaDo00} for lensing reconstruction and by \citet{SmZa06} for general bispectrum analysis. |
for one-point estimator for fy;. | for one-point estimator for $f_{NL}$. |
The analysis depends on finding suitable inverse cosmic variance weighting €C'+ of modes. | The analysis depends on finding suitable inverse cosmic variance weighting $C^{-1}$ of modes. |
It deals with mode-mode coupling in an exact way. | It deals with mode-mode coupling in an exact way. |
In a recent work Munshi&Heavens(2009). further extended this analysis by incorporating two-point statistics or the skew spectrum which we have already introduced above. | In a recent work \citet{MuHe09} further extended this analysis by incorporating two-point statistics or the skew spectrum which we have already introduced above. |
We generalise their results in this work for the case of mixed bispectra for the case of both one-point and two-point studies involving three-way correlations. | We generalise their results in this work for the case of mixed bispectra for the case of both one-point and two-point studies involving three-way correlations. |
The analytical results presented here are being kept as general as possible. | The analytical results presented here are being kept as general as possible. |
However in the next sections we specialise them to individual cases of lensing reconstruction and the mixed bispectrum associated with lensing and the SZ effect as concrete examples. | However in the next sections we specialise them to individual cases of lensing reconstruction and the mixed bispectrum associated with lensing and the SZ effect as concrete examples. |
We ure in constructing an optimal and unbiased estimator for the estimation of mixed skewness (V(O)3(Q0)Z(OQ). | We are interested in constructing an optimal and unbiased estimator for the estimation of mixed skewness $\langle X(\oh)Y(\oh)Z(\oh) \rangle$. |
The fields interestedand similarly for Y and Z. are detined over the entire sky. though observed with a mask and nonuniform noise coverage. | The fields $X(\Omega) =\sum_{lm}X_{lm}Y_{lm}(\oh)$ and similarly for $Y$ and $Z$, are defined over the entire sky, though observed with a mask and nonuniform noise coverage. |
The coverage imprints a mode-mode coupling Cxx]Homil1 in the observed multipoles of a given field in the harmonic space (Xi,Xia | The non-uniform coverage imprints a mode-mode coupling $\left [ C_{XX} \right ]^{-1}_{l_1m_1,l_2m_2}$ in the observed multipoles of a given field in the harmonic space $\langle X_{l_1m_1}X_{l_2m_2}\rangle$. |
For the construction of the optimal estimator it will be useful to define τν as Here Nip represents the harmonics of the data Y withinverse covariance weighting. | For the construction of the optimal estimator it will be useful to define $\tilde X_{l_1m_1}$ as Here $\tilde X_{lm}$ represents the harmonics of the data $X$ with inverse covariance weighting. |
Next we need to deal with the covariance matrix of the modes Nu,4 interms of that of .X;,,,, The auto covariance matrix for VY. C'yx. and that of N.CxXx are related by the following expression: Similarly. the cross-covariance for two different tields Y and Y with inverse variance weighting. in harmonie space can be written as: The estimator that we construct will be based on functions QX.Y.Z] which depends on the input fields. and its derivatives w.r.t. | Next we need to deal with the covariance matrix of the modes $\tilde X_{l_1m_1}$ in terms of that of $X_{l_1m_1}$ The auto covariance matrix for $X$, $C_{XX}$, and that of $\tilde X$, $\tilde C_{XX}$ are related by the following expression: Similarly, the cross-covariance for two different fields $\tilde X$ and $\tilde Y$ with inverse variance weighting, in harmonic space can be written as: The estimator that we construct will be based on functions $\hat Q[\tilde X, \tilde Y, \tilde Z]$ which depends on the input fields, and its derivatives w.r.t. |
the fields e.g. 00X.Y]/0Zi,. | the fields e.g. ${\partial \hat Q[\tilde X, \tilde Y] / \partial \tilde Z_{lm}}$. |
The derivatives are themselves a map with harmonics described by the free indices /in. and are constructed out of two other maps. | The derivatives are themselves a map with harmonics described by the free indices $lm$, and are constructed out of two other maps. |
The function €) on the other hand is an ordinary number which depends on all three input functions and lacks free indices. | The function $\hat Q$ on the other hand is an ordinary number which depends on all three input functions and lacks free indices. |
Similar expressions hold for other fields such as ος 35, Introducing a more compact notation οτε, where oryNoreYoursZ we can write the one-point estimator for the mixed skewness as: This is a main result of the paper. generalising work by Smith&Zaldarriaga(2006) to mixed fields. | Similar expressions hold for other fields such as $\tilde X_{lm}$, $\tilde Y_{lm}$ Introducing a more compact notation $x_i$, where $x_1 = X, x_2 = Y, x_3 = Z$ we can write the one-point estimator for the mixed skewness as: This is a main result of the paper, generalising work by \citet{SmZa06} to mixed fields. |
The ensemble averaging in the linear terms represents Monte-Carlo averaging using simulated non-Gaussian maps. | The ensemble averaging $\langle \rangle$ in the linear terms represents Monte-Carlo averaging using simulated non-Gaussian maps. |
The associated Fisher matrix a scalar in this ease) can be written0 in terms of the functions (2 αι]. its derivative and the cross-covariance matrices involving different fields. | The associated Fisher matrix (a scalar in this case) can be written in terms of the functions $Q[\tilde x_i]$ , its derivative and the cross-covariance matrices involving different fields. |
pointed 1 arcin apart. | pointed 1 arcmin apart. |
This problem was then corrected. resulting in an overall efficiency decrease of the pu. MOSI and Re iustrunents. | This problem was then corrected, resulting in an overall efficiency decrease of the pn, MOS1 and RGS instruments. |
We generated final product using v. 5.L1. | We generated final product using v. 5.4.1. |
A simall soft proton flare event occurred at the bestiis of the observation. we excluded the first 2.5 ks from the analysis. | A small soft proton flare event occurred at the beginning of the observation, we excluded the first 2.5 ks from the analysis. |
We obtained net exposure times of L1 ks and 9 ks for MOS2 aud the other Παςποιές, respectively. | We obtained net exposure times of 14 ks and 9 ks for MOS2 and the other instruments, respectively. |
The source is extremely bright such that the MOS calucras are saturated aud the radial support structure of the mirrors is visible. | The source is extremely bright such that the MOS cameras are saturated and the radial support structure of the mirrors is visible. |
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