source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
. Although the authors quoted their value for 11) as an upper limit. we argue that in principle significant contamination could also be taking place on the ground fine-structure Line. thereby also allecting 11)) and. driving the ratio NOYN in the opposite sense. | Although the authors quoted their value for $^*$ ) as an upper limit, we argue that in principle significant contamination could also be taking place on the ground fine-structure line, thereby also affecting ) and driving the ratio $N^*/N$ in the opposite sense. |
Unfortunately. the ground 1334 line is often heavily saturated: to circunvent this problem there have been many alternative approaches to derive the 11) column density bv other indirect methods. | Unfortunately, the ground 1334 line is often heavily saturated; to circunvent this problem there have been many alternative approaches to derive the ) column density by other indirect methods. |
Prochaska (1999). used the ratio of NCCIDD/N(GeO12) in à velocity region where the σοι line was not saturated to derive the corresponding value at the component where the line was detected. | Prochaska \shortcite{P99} used the ratio of ) in a velocity region where the ground line was not saturated to derive the corresponding value at the component where the line was detected. |
Outram. Challee Carswell (1999) assumed a carbon abundance relative to iron. Fe} >-0.3 to obtain a tighter lower limit on the 1) column density in a DLA svstem at 2,1,=2.62 towards GD1759τὸ, | Outram, Chaffee Carswell \shortcite{OCC} assumed a carbon abundance relative to iron $>$ -0.3 to obtain a tighter lower limit on the ) column density in a DLA system at $z_{\rmn{abs}}=2.62$ towards GB1759+75. |
In our sample we have included only direct. measurements on the column densities. | In our sample we have included only direct measurements on the column densities. |
In sections 3.1--3.2. below. we will separately study the DLA and LL systems in our sample. | In sections \ref{section:DLA}- \ref{section:LL} below, we will separately study the DLA and LL systems in our sample. |
Again. as a working hypothesis we shall assume the temperature-redshift relation as predicted. by the standard model. | Again, as a working hypothesis we shall assume the temperature-redshift relation as predicted by the standard model. |
Ehe. validity of this relation is discussed in section 3.3.. | The validity of this relation is discussed in section \ref{section:CMBR}. |
DLA svstenis have very high. neutral hydrogen column densities (logN(LLEI) 20.3). | DLA systems have very high neutral hydrogen column densities $\log\rmn{N}(\hbox{H\,{\sc i}})>20.3$ ). |
This makes them elfectively shiclelecl from the ionizing radiation. causing their contents o be essentially neutral materia (Viegas 1995).. | This makes them effectively shielded from the ionizing radiation, causing their contents to be essentially neutral material \cite{Viegas95}. . |
We use the fine-structure lines column density. ratios observed in the DLA systems listed in table 2. to set upper imits to their neutral hydrogen volume densities mye ancl to he intensities of the UV. radiation field present. | We use the fine-structure lines column density ratios observed in the DLA systems listed in table \ref{obsdata} to set upper limits to their neutral hydrogen volume densities $n_{\rmn{H}^0}$ and to the intensities of the UV radiation field present. |
Given the ugh neutral hydrogen. column density. probably all of the ivdrogen ionizing radiation will be absorbed. leaving very ew photons with energies greater than 1 Rwd. | Given the high neutral hydrogen column density, probably all of the hydrogen ionizing radiation will be absorbed, leaving very few photons with energies greater than 1 Ryd. |
The spectral shape of the UV radiation field willthen be similar to the one found in our own galaxy. and we therefore assume the | The spectral shape of the UV radiation field willthen be similar to the one found in our own galaxy, and we therefore assume the |
Figures | and 2. reveal a smooth evolution along the sequence ΚΗΙ — MIL — Ba/S. in the sense that the upper boundary of the populated region in the («logP) diagram moves towards longer periods (this ts reflected by the three curved lines which roughly delineate the regions populated by these three classes. their exact definition being given below). | Figures \ref{Fig:elogP_M} and \ref{Fig:elogP_panels} reveal a smooth evolution along the sequence KIII – MIII – Ba/S, in the sense that the upper boundary of the populated region in the $(e -
\log P)$ diagram moves towards longer periods (this is reflected by the three curved lines which roughly delineate the regions populated by these three classes, their exact definition being given below). |
This is clearly a consequence of the larger radii reached by stars evolving along this sequence. | This is clearly a consequence of the larger radii reached by stars evolving along this sequence. |
In the case of Ba and Te-poor S giants. 1t is actually their white dwarf (WD) companions which reached very large radii while evolving on the AGB. | In the case of Ba and Tc-poor S giants, it is actually their white dwarf (WD) companions which reached very large radii while evolving on the AGB. |
For K giants. the situation is in principle somewhat more complex. since this class mixes stars on the first giant branch and stars in the core He-burning phase. | For K giants, the situation is in principle somewhat more complex, since this class mixes stars on the first giant branch and stars in the core He-burning phase. |
stars belonging to the latter category have gone through the RGB tip. where they reached a very large radius (similar to. or even larger than that of M giants). | stars belonging to the latter category have gone through the RGB tip, where they reached a very large radius (similar to, or even larger than that of M giants). |
Therefore. if those low-mass. core-He burning stars were to dominate among K giants. their distribution in the (c.logP) diagram should be characterised by an envelope located at even longer periods than that for M giants. | Therefore, if those low-mass, core-He burning stars were to dominate among K giants, their distribution in the $(e - \log P)$ diagram should be characterised by an envelope located at even longer periods than that for M giants. |
Fig. | Fig. |
shows that this is not the case. because the sample of open- K giants plotted in Fig. | \ref{Fig:elogP_panels} shows that this is not the case, because the sample of open-cluster K giants plotted in Fig. |
2. ts in fact dominated by intermediate-mass stars. as may be judged from the turnoff masses of the corresponding clusters. most of them being larger than 2 citepMermilliod-2007b.. | \ref{Fig:elogP_panels} is in fact dominated by intermediate-mass stars, as may be judged from the turnoff masses of the corresponding clusters, most of them being larger than 2 \\citep{Mermilliod-2007b}. |
The complication introduced by the mixture of evolutionary states among K giants is thus not a concern. | The complication introduced by the mixture of evolutionary states among K giants is thus not a concern. |
Equating the stellar radius to the Roche radius results in a threshold period (for given component masses) below which the primary star undergoes RLOF. | Equating the stellar radius to the Roche radius results in a threshold period (for given component masses) below which the primary star undergoes RLOF. |
Adopting Paezynisski’s usual expression for the Roche radius Πρ around star | where g=AL,/ALS and A is the orbital separation. one finds that a star of radius 40 R.. fills its Roche lobe in a system of period P=το d. for masses M4=1.3 M. and AL,—0.6M ... | Adopting Paczyńsski's usual expression for the Roche radius $R_{R,1}$ around star 1 where $q = M_1/M_2$ and $A$ is the orbital separation, one finds that a star of radius 40 $_\odot$ fills its Roche lobe in a system of period $P = 70$ d, for masses $M_1 = 1.3 $ $_\odot$ and $M_2 =
0.6$ $_\odot$. |
Although the Roche lobe concept is in principle only applicable to circular orbits. one may formally compute the orbital periods for which the primary star fills its Roche lobe periastron. by replacing A by A(1ο) in the above expression. | Although the Roche lobe concept is in principle only applicable to circular orbits, one may formally compute the orbital periods for which the primary star fills its Roche lobe , by replacing $A$ by $A(1-e)$ in the above expression. |
It is quite remarkable that the relationship between P and e so obtained (assuming Ry=36 .) exactly matches the boundary of the region occupied by KIII giants in the (6.logP) diagram. both for cluster and giants (Fig. 2)). | It is quite remarkable that the relationship between $P$ and $e$ so obtained (assuming $R_R = 36$ $_\odot$ ) exactly matches the boundary of the region occupied by KIII giants in the $(e - \log P)$ diagram, both for cluster and giants (Fig. \ref{Fig:elogP_panels}) ). |
This excellent match thus clearly suggests that mass transfer at periastron plays a crucial role in shaping the («logP?) diagram (?).. | This excellent match thus clearly suggests that mass transfer at periastron plays a crucial role in shaping the $(e - \log P)$ diagram \citep{Soker00}. |
It may seem surprising that the “periastron envelope’ LA©) 2 constant. or 1€) = constant] represents a better fit to the data than the ‘circularisation envelope’ [ACLc2) = constant. or P7/7(1.—£7) = constant: see Fig. | It may seem surprising that the 'periastron envelope' $A(1-e)$ = constant, or $P^{2/3}(1-e)$ = constant] represents a better fit to the data than the 'circularisation envelope' $A(1-e^2)$ = constant, or $P^{2/3}(1-e^2)$ = constant; see Fig. |
6 of Paper II]. resulting from the fact that circularisation keeps the angular momentum per unit reduced mass constant (???)., | 6 of Paper II], resulting from the fact that circularisation keeps the angular momentum per unit reduced mass constant \citep{Zahn-1977,Hut81,Duquennoy-92}. |
Indeed. às the star gets closer to its Roche lobe. it should circularise first and then fill its Roche lobe. and possibly disappear from the sample due to cataclysmic mass transfer. | Indeed, as the star gets closer to its Roche lobe, it should circularise first and then fill its Roche lobe, and possibly disappear from the sample due to cataclysmic mass transfer. |
The samples of K-giant binaries clearly favour the periastron envelope over thecircularisation envelope. | The samples of K-giant binaries clearly favour the periastron envelope over thecircularisation envelope. |
The reason for this may be the following. | The reason for this may be the following. |
When à system is close to filling its | When a system is close to filling its |
One system, CXGG0095951+0140.8,has Amij.=—-0.014mag (i.e., Ami»=2.10X-0.02mag as obtained from the difference between the magnitudes in the rest-frame R200=832+19.6 M200=9.5(+0.42)xr-band),1013 Mo, z=0.372, and six kpc,spectroscopic members. | One system, $+$ 0140.8,has $\mathrm{\Delta m_i}_{12} = 2.19 \pm 0.014~\mathrm{mag}$ (i.e., $\mathrm{\Delta m}_{12} = 2.10 \pm 0.02~\mathrm{mag}$ as obtained from the difference between the magnitudes in the rest-frame $r$ -band), $R_{200} = 832 \pm 19.6~\mathrm{kpc}$, $M_{200} = 9.5~(\pm 0.42) \times 10^{13}~\mathrm{M}_{\sun}$ , $z = 0.372$, and six spectroscopic members. |
It is part of a large-scale structure (see fig. | It is part of a large-scale structure (see fig. |
1 in (09 and fig. | 1 in G09 and fig. |
3 in Scoville et al. | 3 in Scoville et al. |
2007) populated by 28 X-ray emitting groups distributed across the entire 2deg? area of the (corresponding to a cross size of about 25.5Mpc at z= 0.37). | 2007) populated by 28 X-ray emitting groups distributed across the entire $2~\mathrm{deg}^2$ area of the (corresponding to a cross size of about $25.5~\mathrm{Mpc}$ at $z=0.37$ ). |
Furthermore, its dominant galaxy (with rest-frame absolute magnitude M;= —24.87) hosts apoint-likea radio source (see Giodini et al. | Furthermore, its dominant galaxy (with a rest-frame absolute magnitude $M_{i} = -24.87$ ) hosts apoint-like radio source (see Giodini et al. |
2010). | 2010). |
The other fossil group, 0095951+0212.6, has Ami,=2.35+0.014mag (ie., Ami»=2.32+0.02mag in the rest-frame r-band), Haeo=478+54.4kpc, Μου=1.9(£0.41)x10?Mo, z=0.425, and eight spectroscopic members. | The other fossil group, $+$ 0212.6, has $\mathrm{\Delta m_i}_{12} = 2.35 \pm 0.014~\mathrm{mag}$ (i.e., $\mathrm{\Delta m}_{12} = 2.32 \pm 0.02~\mathrm{mag}$ in the rest-frame $r$ -band), $R_{200} = 478 \pm 54.4~\mathrm{kpc}$, $M_{200} = 1.9~(\pm 0.41) \times 10^{13}~\mathrm{M}_{\sun}$, $z = 0.425$, and eight spectroscopic members. |
It is isolated and its BCG has M;=—23.87. | It is isolated and its BCG has $M_{i} = -23.87$. |
Basic properties of the two fossil groups under study are listed in Table 1. | Basic properties of the two fossil groups under study are listed in Table 1. |
In addition, we note that they populate the upper half of the distribution of the X-ray selected groups in the (galaxy) stellar mass fraction-group total-mass diagram (G09, their fig. | In addition, we note that they populate the upper half of the distribution of the X-ray selected groups in the (galaxy) stellar mass fraction–group total-mass diagram (G09, their fig. |
5), where quantities are estimated at 0.7 R2oo. | 5), where quantities are estimated at $0.7 R_{200}$ . |
This is particularly true for 0095951--0212.6. | This is particularly true for $+$ 0212.6. |
As from GO095,, the stellar mass | As from the stellar mass |
Zuckermanetal.L972:: Sclilkeetal.1992:: Ilrotaοal.1998.. C. HCNII!. al.2002)). ITCNIT! ΝΠΗΟΞΠΝΟ Herbst.Terzieva.&Talbi2000)) was bv Irvinectal.(1996) in comet 61996 D2 (νακακο), | \citealt{snyder72, zucker72}; \citealt{black76}) \citealt{schilke92}; \citealt{hirota98}, $^+$ $^+$ \citealt{rodgers01a, charnley02}) $^+$ $_2$ \citealt{herbst00}) was by \cite{irvine96} in comet C/1996 B2 (Hyakutake). |
The measured INCΠο abundance ratio..L. was simular to that iu iuterstellar clouds with ac9 temperatureOT of order 50 Ix. sugecstinge that coluctary TNC may be unprocessed interstellar material incorporated inte the comets nucleus. | The measured HNC/HCN abundance ratio, was similar to that in interstellar clouds with a temperature of order 50 K, suggesting that cometary HNC may be unprocessed interstellar material incorporated into the comet's nucleus. |
IHToscever. levinectal.(1996) argued that a nuuber of alternative processes may also explain the observed INC/IICN ratio iu comet IIvakutake. including irradiation of icy matrix containiug IICN. non-equilibrium chemical processes in the solar nebula. gas-phase processes iu the coma itself. infrared relaxation of ICN from excited vibrational levels of the eround electronic state. or photo-dissociation of a heavier pareut molecule. | However, \cite{irvine96} argued that a number of alternative processes may also explain the observed HNC/HCN ratio in comet Hyakutake, including irradiation of icy matrix containing HCN, non-equilibrium chemical processes in the solar nebula, gas-phase processes in the coma itself, infrared relaxation of HCN from excited vibrational levels of the ground electronic state, or photo-dissociation of a heavier parent molecule. |
A stroug variation of the IENC/IICN abundauce ratio πι comet C/1995 OL (ILde-Bopp) with heliocentric distance (from at 2.9 AU to πο 1 AU: Biverctal.1997: Lhwineetal. 1998)) questioned he interstellar origin of cometary TNC and sugeestedCoco a production imechanisu du the coma itself as a nore likely explanation. | A strong variation of the HNC/HCN abundance ratio in comet C/1995 O1 (Hale-Bopp) with heliocentric distance (from at 2.9 AU to near 1 AU; \citealt{biver97}; \citealt{irvine98}) ) questioned the interstellar origin of cometary HNC and suggested a production mechanism in the coma itself as a more likely explanation. |
Rodgers&Charuley(1998) srescuted a comprehensive model of the cometary coma chemistry and suggested that iu very active comets. such as comet ILale-Dopp. the observed variation of he TNC abundance with the helioceutrie distance. as observed with sinele-dish telescopes. cau be explained wojsolerization of IIC'N driven by the iupact of ast hwdroseu atoms produced in photo-dissociatiou of went molecules. | \cite{rodgers98} presented a comprehensive model of the cometary coma chemistry and suggested that in very active comets, such as comet Hale-Bopp, the observed variation of the HNC abundance with the heliocentric distance, as observed with single-dish telescopes, can be explained by isomerization of HCN driven by the impact of fast hydrogen atoms produced in photo-dissociation of parent molecules. |
However. their model overproduced he TING abundance at ~3 AU by abot a factor of 2. | However, their model overproduced the HNC abundance at 3 AU by about a factor of 2. |
Ina subsequent paper Rodgers&Charuley(2001) showed that the same imechauign cannot reproduce observed IENC'/IICN. abundance ratios iu nmoderatelv active comoets at —1 AU. such as cometC7/1999 ITI (Loo). | In a subsequent paper \cite{rodgers01b} showed that the same mechanism cannot reproduce observed HNC/HCN abundance ratios in moderately active comets at 1 AU, such as cometC/1999 H1 (Lee). |
The applicability of the model to very active comets has also been questioned by interferometric observations of | The applicability of the model to very active comets has also been questioned by interferometric observations of |
Of the 13 departures from tolerable fits to redshifts and distances, eight belong to four misfit galaxies, Cetus, Tucana, DDO 210, and the Sagittarius dwarf irregular. | Of the 13 departures from tolerable fits to redshifts and distances, eight belong to four misfit galaxies, Cetus, Tucana, DDO 210, and the Sagittarius dwarf irregular. |
They have similar discrepancies in redshifts and distances, and they have similar orbits (plotted as the dashed lines in Figs. | They have similar discrepancies in redshifts and distances, and they have similar orbits (plotted as the dashed lines in Figs. |
3 to 5)) that emanate from SGL~200°, SGB~30°. | \ref{Fig:3} to \ref{Fig:5}) ) that emanate from $SGL\sim 200^\circ$, $SGB\sim 30^\circ$. |
It may be significant that this is in the direction of the Local Void (Tully et al. | It may be significant that this is in the direction of the Local Void (Tully et al. |
2008). | 2008). |
The common features — low redshifts, large distances, and similar orbits — argue against the idea that measurement errors are to blame, and invites the speculation that they are manifestations of an inadequate external mass model that would have to be particularly serious near these four galaxies. | The common features — low redshifts, large distances, and similar orbits — argue against the idea that measurement errors are to blame, and invites the speculation that they are manifestations of an inadequate external mass model that would have to be particularly serious near these four galaxies. |
If an adjustment of our phenomenological representation of the external mass could reduce the peculiar gravitational acceleration toward MW near the four misfits it would allow larger redshifts at lower present distances, in the direction wanted to improve the fit. | If an adjustment of our phenomenological representation of the external mass could reduce the peculiar gravitational acceleration toward MW near the four misfits it would allow larger redshifts at lower present distances, in the direction wanted to improve the fit. |
A search for a fifth external mass capable of producing this effect has not yielded anything promising, however. | A search for a fifth external mass capable of producing this effect has not yielded anything promising, however. |
An explanation of the enigmatic properties of these four misfit galaxies remains an interesting open issue. | An explanation of the enigmatic properties of these four misfit galaxies remains an interesting open issue. |
(EW) baseline vectors. | (EW) baseline vectors. |
Next. we cdeseribe the. estimation of array geometry. | Next, we describe the re-estimation of array geometry. |
We begin with a brief description of the mode of observations with MICE. | We begin with a brief description of the mode of observations with MRT. |
MICE has 32 fixed antennas in the EW arm and 15 movable antenna trollevs in the NS arm. | MRT has 32 fixed antennas in the EW arm and 15 movable antenna trolleys in the NS arm. |
For measuring visibilities. the 15 NS trollevs are configured by spreading them over S4 m with an inter-trollev spacing of 6 m (to avoid shadowing of one trolley by another). | For measuring visibilities, the 15 NS trolleys are configured by spreading them over 84 m with an inter-trolley spacing of 6 m (to avoid shadowing of one trolley by another). |
MICE measures cdillerent. Fourier components of the. brightness clistribution of the sky in 63 dillerent configurations (referred lo as allocations) to. sample NS baselines every mum. Therefore. ellectively. there are 945 antenna positions (63 allocations * 15 antennas/allocation) in the NS arm and a total of 30.240 (945 * 32) visibilities are used for imagine. | MRT measures different Fourier components of the brightness distribution of the sky in 63 different configurations (referred to as ) to sample NS baselines every m. Therefore, effectively, there are 945 antenna positions (63 allocations * 15 antennas/allocation) in the NS arm and a total of 30,240 (945 * 32) visibilities are used for imaging. |
A small error in a measuring scale of relatively shorter leneth is likely to build up systematically while establishing the geometry of longer basclines. | A small error in a measuring scale of relatively shorter length is likely to build up systematically while establishing the geometry of longer baselines. |
This οσο would. be observed. in the instrumental. phases estimated using different calibrators. | This effect would be observed in the instrumental phases estimated using different calibrators. |
In. principle. the instrumental phases estimated using two calibrators at dilferent. declinations. for a given baseline. should be the same. allowing for temporal variations in the instrumental gains. | In principle, the instrumental phases estimated using two calibrators at different declinations, for a given baseline, should be the same, allowing for temporal variations in the instrumental gains. |
A non-zero cdillerence in these estimates may be due to positional errors of the baseline or positions of calibrators. | A non-zero difference in these estimates may be due to positional errors of the baseline or positions of calibrators. |
As mentioned earlier. our analysis of positional error in sources and the homography matrix cued to positional errors in. baselines (or antenna positions). | As mentioned earlier, our analysis of positional error in sources and the homography matrix cued to positional errors in baselines (or antenna positions). |
The simple principle of astrometry CEhomsonetal.2001) was used to estimate errors in antenna. positions and is discussed below. | The simple principle of astrometry \citep{book:thomson} was used to estimate errors in antenna positions and is discussed below. |
The observed. visibility phase. (5/. in a baseline with components (u,;;.0;05). due to calibrator S, with direction cosines (f°!mS!nF!) is given by: Where.⇁ ορETTui represents (rue instrumental; phases. ὁ;=1.2.....32 represents EW antennas and j=1.2.....045 represents NS antennas. | The observed visibility phase, $\psi_{ij}^{\mathcal{S}^{}_1}$, in a baseline with components $\left(u^{}_{ij}, v^{}_{ij}, w^{}_{ij}\right)$, due to calibrator $S^{}_{1}$ with direction cosines $\left(l^{\mathcal{S}_1},m^{\mathcal{S}_1},n^{\mathcal{S}_1}\right)$, is given by: Where, $\phi_{ij}^{\mbox{\small ins}}$ represents true instrumental phases, $i=1,2,\ldots,32$ represents EW antennas and $j=1,2,\ldots,945$ represents NS antennas. |
For moeridian transit imaging Equation 7 becomes: The instrumental phases; 657.δι estimated. using. the measured geometry are given by: Here. Ae; and Aw;, are errors in the assumed. baseline | For meridian transit imaging Equation \ref{e:obsphasebasiceqn} becomes: The instrumental phases, $\phi_{ij}^{\mathcal{S}^{}_{1}}$, estimated using the measured geometry are given by: Here, $\Delta v^{}_{ij}$ and $\Delta w^{}_{ij}$ are errors in the assumed baseline |
center to the current position of the bubble and M(R) ts the total gravitating mass within R. then where we have used AR=34" (2.6kpe). the projected distance from the cluster center. and M(R)=1.4«10!!M. (Cótté et 22001). | center to the current position of the bubble and $M(R)$ is the total gravitating mass within $R$, then where we have used $R = 34''$ $2.6\,{\rm kpc}$ ), the projected distance from the cluster center, and $M(R) = 1.4\times
10^{11}\,M_{\odot}$ (Côtté et 2001). |
Cy~0.5 is the drag coefficient. for a roughly spherical bubble. | $C_{W} \sim 0.5$ is the drag coefficient for a roughly spherical bubble. |
Since the actual distance to the cluster center almost certainly exceeds the projected distance. this gives a lower limit for the rise time. | Since the actual distance to the cluster center almost certainly exceeds the projected distance, this gives a lower limit for the rise time. |
Furthermore. according to (1). the speed of the bubble. 383kms!. exceeds half of the sound speed and so is overestimated. | Furthermore, according to \ref{eq:risetime}) ), the speed of the bubble, $383\rm\ km\ s^{-1}$, exceeds half of the sound speed and so is overestimated. |
Thus its rise time is underestimated. even if the budding bubble les in the plane of the sky. | Thus its rise time is underestimated, even if the budding bubble lies in the plane of the sky. |
During its rapid mitial expansion. the boundary of the bubble will generally be stable. | During its rapid initial expansion, the boundary of the bubble will generally be stable. |
As a result. the motion of the bubble boundary generally needs to be subsonic before a bubble even starts to form. | As a result, the motion of the bubble boundary generally needs to be subsonic before a bubble even starts to form. |
This adds a further delay to μις after the outburst. but before the bubble is formed. | This adds a further delay to $\tau_{\rm bubble}$ after the outburst, but before the bubble is formed. |
If we do associate the budding bubble with an energetic nuclear event. then the constraints on its formation timescale make it quite reasonable to associate it with the current outburst (associated with the jet) that commenced about 107 years ago. | If we do associate the budding bubble with an energetic nuclear event, then the constraints on its formation timescale make it quite reasonable to associate it with the current outburst (associated with the jet) that commenced about $10^7$ years ago. |
The most striking X-ray features in M87 are the two arms that extend east and southwest from the inner lobe region. | The most striking X-ray features in M87 are the two arms that extend east and southwest from the inner lobe region. |
These also are seen in the 90 em image (see Fig. | These also are seen in the 90 cm image (see Fig. |
11. fora composite X-ray-radio view of M87). | \ref{fig:overlay} for a composite X-ray-radio view of M87). |
Previous spectroscopic studies of the arms have utilized the XMM-Newton observations (Belsole et al. | Previous spectroscopic studies of the arms have utilized the XMM-Newton observations (Belsole et al. |
2001. Molendi 2002). | 2001, Molendi 2002). |
They find that the arms are cool and portions are poorly fit by single temperature components. | They find that the arms are cool and portions are poorly fit by single temperature components. |
Our Chandra results agree with these previous analyses. as does the XMM-Newton temperature map (Fig. 6)). | Our Chandra results agree with these previous analyses, as does the XMM-Newton temperature map (Fig. \ref{fig:xmm_tmap}) ). |
We find that the arms require at least two components (with variable abundances. VMEKAL or VAPEC) with the low and high temperature components in the range 1-1.5 keV and 2-2.7 keV respectively. | We find that the arms require at least two components (with variable abundances, VMEKAL or VAPEC) with the low and high temperature components in the range 1-1.5 keV and 2-2.7 keV respectively. |
Although the two arms are likely related to the same outburst. we discuss each separately. | Although the two arms are likely related to the same outburst, we discuss each separately. |
The eastern X-ray and radio arm begins at the eastern edge of the inner radio cocoon. but its appearance is much more amorphous than that of the southwestern arm (see Fig. [.. 3.. | The eastern X-ray and radio arm begins at the eastern edge of the inner radio cocoon, but its appearance is much more amorphous than that of the southwestern arm (see Fig. \ref{fig:bl1sum}, \ref{fig:flatbl1_sm2}, |
and 4)). | and \ref{fig:divking}) ). |
At the base of the filament (Fig. | At the base of the filament (Fig. |
Subsets and Splits
No saved queries yet
Save your SQL queries to embed, download, and access them later. Queries will appear here once saved.