source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
1. and 2)) are at least four bubbles with sizes comparable to that of the "bud" discussed above and streamers of gas bounding these buoyantly rising bubbles. | \ref{fig:bl1sum} and \ref{fig:adapt}) ) are at least four bubbles with sizes comparable to that of the “bud” discussed above and streamers of gas bounding these buoyantly rising bubbles. |
Typical bubble sizes are ~10” (0.8 kpe) in radius and are reminiscent of the “effervescent” heating described by Begelman (2003). | Typical bubble sizes are $\sim10''$ (0.8 kpc) in radius and are reminiscent of the “effervescent” heating described by Begelman (2003). |
Fig. | Fig. |
9 shows a projection across one of these “effervescent” bubbles 1.25" (5.8 kpe) east of the M87 nucleus (labeled "bubble" in Fig. | \ref{fig:bubble_proj} shows a projection across one of these “effervescent” bubbles $1.25'$ (5.8 kpc) east of the M87 nucleus (labeled “bubble” in Fig. |
lee). | \ref{fig:bl1sum}c c). |
The temperature structure (Fig. 6)) | The temperature structure (Fig. \ref{fig:xmm_tmap}) ) |
of the eastern arm shows X-ray features that are consistent with cool matertal uplifted by a rising torus (Churazov et al. | of the eastern arm shows X-ray features that are consistent with cool material uplifted by a rising torus (Churazov et al. |
2001). | 2001). |
First. the largest concentration of the coolest gas lies midway along the eastern arm (1/—2 from M87's nucleus). | First, the largest concentration of the coolest gas lies midway along the eastern arm $1'-2'$ from M87's nucleus). |
Second. the cool gas column in the eastern arm narrows at the edge of the radio torus closest to the M87 nucleus and then broadens within the torus (labeled "Uplifted Gas" in Fig. | Second, the cool gas column in the eastern arm narrows at the edge of the radio torus closest to the M87 nucleus and then broadens within the torus (labeled “Uplifted Gas” in Fig. |
4bb). just as one might expect for gas uplifted by a buoyant toroidal plasma bubble (see Fig. | \ref{fig:divking}b b), just as one might expect for gas uplifted by a buoyant toroidal plasma bubble (see Fig. |
11. and Fig. | \ref{fig:overlay} and Fig. |
3 and 4 in Churazov et al. | 3 and 4 in Churazov et al. |
2001). | 2001). |
A projection along the arm. Fig. 10.. | A projection along the arm, Fig. \ref{fig:eastern_arm_proj}, |
shows a brightening at the radial distance of the 14 kpe ring. | shows a brightening at the radial distance of the 14 kpc ring. |
A similar brightening occurs at about the same angular distance on the southwestern arm. | A similar brightening occurs at about the same angular distance on the southwestern arm. |
While the feature in the southwestern arm is partially obscured by the change from the ACIS S3 to S2 chip in the Chandra image. it is clearly seen in both the ROSAT HRI and XMM-Newton images (Fig. 5)). | While the feature in the southwestern arm is partially obscured by the change from the ACIS S3 to S2 chip in the Chandra image, it is clearly seen in both the ROSAT HRI and XMM-Newton images (Fig. \ref{fig:rosat}) ). |
[f this brightening ts associated with the passage of the same shock that produced the ring. then this arm (and the southwestern arm as well) must lie close to the plane of the sky. | If this brightening is associated with the passage of the same shock that produced the ring, then this arm (and the southwestern arm as well) must lie close to the plane of the sky. |
If this brightening does arise from the passage of the shock. it is likely that the so called "radio ear". the vortex- structure that forms the end of the bright eastern radio filament (see Fig. 11)). | If this brightening does arise from the passage of the shock, it is likely that the so called “radio ear”, the vortex-like structure that forms the end of the bright eastern radio filament (see Fig. \ref{fig:overlay}) ), |
falls between the shocks associated with the 14 kpe and 17 kpe rings. | falls between the shocks associated with the 14 kpc and 17 kpc rings. |
This could alternatively explain the flat. ring-like appearance of this radio feature. since passage of a shock through a bubble of relativistic plasma embedded in a background of cold thermal material will induce strong vorticity in the plasma. turning it into a ring-like structure (Ensslin Bruggen 2002). | This could alternatively explain the flat, ring-like appearance of this radio feature, since passage of a shock through a bubble of relativistic plasma embedded in a background of cold thermal material will induce strong vorticity in the plasma, turning it into a ring-like structure (Ensslin Bruggen 2002). |
Combined with the effect of vorticity creation in buoyantly rising bubbles described by Churazov et al. ( | Combined with the effect of vorticity creation in buoyantly rising bubbles described by Churazov et al. ( |
2001). this could account for the rather filamentary appearance of this feature. | 2001), this could account for the rather filamentary appearance of this feature. |
At the end of the eastern arm (~3! east of the M87 nucleus). the X-ray image (Fig. 4)) | At the end of the eastern arm $\sim3'$ east of the M87 nucleus), the X-ray image (Fig. \ref{fig:divking}) ) |
shows an almost circular enhancement (radius of I’ centered at RA=12:31:05.397 DEC=+12:25:10.01) extending to the north (beyond the northern "ear" of the radio emitting torus). | shows an almost circular enhancement (radius of $1'$ centered at RA=12:31:05.397 DEC=+12:25:10.01) extending to the north (beyond the northern “ear” of the radio emitting torus). |
This circular feature is bounded on three sides by X-ray enhancements (see Fig. 4) | This circular feature is bounded on three sides by X-ray enhancements (see Fig. \ref{fig:divking}) ) |
which originate at the eastern arm and it is bounded to the northwest by a pair of radio ares (best seen in the 90 cm image: see Fig. 11). | which originate at the eastern arm and it is bounded to the northwest by a pair of radio arcs (best seen in the 90 cm image; see Fig. \ref{fig:overlay}) ). |
The X-ray temperature of this circular region is intermediate in temperature (1.8-1.9 keV) as seen in Fig. | The X-ray temperature of this circular region is intermediate in temperature (1.8-1.9 keV) as seen in Fig. |
6 and is comparable to that of the end of southeastern arm (as it swings to the east). | \ref{fig:xmm_tmap} and is comparable to that of the end of southeastern arm (as it swings to the east). |
The two enhancements. labeled andE2 in Fig. | The two enhancements, labeled and in Fig. |
dec). which bound the circular region. appear similar to the two filaments into which the southwestern arm divides (see below). | \ref{fig:divking}c c), which bound the circular region, appear similar to the two filaments into which the southwestern arm divides (see below). |
We suggest that the outer portions of the eastern arm are similar to the southwestern arm. but seen from a different orientation. | We suggest that the outer portions of the eastern arm are similar to the southwestern arm, but seen from a different orientation. |
The southwestern X-ray arm originates (see Fig. | The southwestern X-ray arm originates (see Fig. |
| and Fig. 2)) | \ref{fig:inner_cocoon}
and Fig. \ref{fig:adapt}) ) |
as a narrow filament of width approximately 10” (0.8 kpc) at its narrowest when it exits from the bright inner core (at a distance of 50”. 3.9 kpe from the nucleus). | as a narrow filament of width approximately $10''$ (0.8 kpc) at its narrowest when it exits from the bright inner core (at a distance of $50''$, 3.9 kpc from the nucleus). |
The filament extends in an almost straight line to the southwest for ~2’ (9.3 kpe). | The filament extends in an almost straight line to the southwest for $\sim 2'$ (9.3 kpc). |
As seen in Fig. 11.. | As seen in Fig. \ref{fig:overlay}, |
over this distance it appears uncorrelated with the radio filament that extends in approximately the same direction. | over this distance it appears uncorrelated with the radio filament that extends in approximately the same direction. |
At a distance of about 3.4 (15.8 kpe). the X-ray filament bifurcates (the two sections are labeled and in Fig. | At a distance of about $3.4'$ (15.8 kpc), the X-ray filament bifurcates (the two sections are labeled and in Fig. |
4ec). | \ref{fig:divking}c c). |
and the correspondence between the radio plasma and X-ray gas becomes more direct. | and the correspondence between the radio plasma and X-ray gas becomes more direct. |
The brightest radio emission lies between the two X-ray arms as they both rotate clockwise in the plane of the sky and eventually turn due east. | The brightest radio emission lies between the two X-ray arms as they both rotate clockwise in the plane of the sky and eventually turn due east. |
Young et al ( | Young et al. ( |
2002) suggested that the arms are overpressurized. | 2002) suggested that the arms are overpressurized. |
Assuming the southwestern arm is a cylinder lying in the plane of the sky. we find that the pressure in the arm is roughly twice that of the hotter ambient gas. | Assuming the southwestern arm is a cylinder lying in the plane of the sky, we find that the pressure in the arm is roughly twice that of the hotter ambient gas. |
We found no elemental abundance differences that could explain | We found no elemental abundance differences that could explain |
In the local universe extraplanar gas is detected in most highly inclined galaxies that have total infrared luminosities >3x10Lo | In the local universe extraplanar gas is detected in most highly inclined galaxies that have total infrared luminosities $>3\times 10^{10}\ \lsol$. |
Our IFU observations include three such super-group 2005).galaxies and we discover that two have extraplanar emission with Ha and [ΝΤ FWHM line-widths of 50—150kms! (Fig. 3)). | Our IFU observations include three such super-group galaxies and we discover that two have extraplanar emission with $\alpha$ and [NII] FWHM line-widths of $50-150\kms$ (Fig. \ref{fig:s23}) ). |
$G1120-82 is a disk-dominated member viewed nearly edge-on that lies on the infrared-radio relation for local star-forming galaxies and is not detected with | SG1120-S2 is a disk-dominated member viewed nearly edge-on that lies on the infrared-radio relation for local star-forming galaxies and is not detected with (Table \ref{tab:sour}) ). |
We detect [NII] and Ha emission in the disk (Tableand, 1)).surprisingly, also at projected heights of ry~7.5 kpc above the disk (Fig. 3)) | We detect [NII] and $\alpha$ emission in the disk and, surprisingly, also at projected heights of $r_{h}\sim7.5$ kpc above the disk (Fig. \ref{fig:s23}) ). |
In the spaxels sampling the disk, the ratios of 0 to 0.2 are consistent with shocked gas and log[NII]/Hathe measured line-widths correspond to gas velocities of ~300—400km s!. | In the spaxels sampling the disk, the $\log$ $\alpha$ ratios of 0 to 0.2 are consistent with shocked gas and the measured line-widths correspond to gas velocities of $\sim300-400\kms$ . |
In the extraplanar spaxels (τι>5 kpc), the log[NII|/Ha ratios of -0.2 to -0.4 are consistent with photoionization by starlight and the line-widths correspond to gas velocities of ~50—150kmsο. | In the extraplanar spaxels $r_{h}>5$ kpc), the $\log$ $\alpha$ ratios of -0.2 to -0.4 are consistent with photoionization by starlight and the line-widths correspond to gas velocities of $\sim50-150~\kms$. |
Our multi-wavelength observations indicate that as with SG1120-S1, the shocked gas in the central region of $G1120-82 is due to star formation and not an AGN. | Our multi-wavelength observations indicate that as with SG1120-S1, the shocked gas in the central region of SG1120-S2 is due to star formation and not an AGN. |
$G1120-S3 is also an inclined disk-dominated member with comparable IR luminosity to SG1120-S2 (Table 1)); it is not detected in the radio nor X-ray observations. | SG1120-S3 is also an inclined disk-dominated member with comparable IR luminosity to SG1120-S2 (Table \ref{tab:sour}) ); it is not detected in the radio nor X-ray observations. |
The IFU maps show and Ha emission in both the disk and extraplanar spaxels[NIJ] (rj~10 kpc), and the line- are consistent with photoionization by starlight. | The IFU maps show [NII] and $\alpha$ emission in both the disk and extraplanar spaxels $r_{h}\sim10$ kpc), and the line-ratios are consistent with photoionization by starlight. |
The FWHM line-widths correspond to velocities of ~200—300kms! iin the disk spaxels and decrease to —50150kms! above the disk. | The FWHM line-widths correspond to velocities of $\sim200-300~\kms$ in the disk spaxels and decrease to $\sim50-150~\kms$ above the disk. |
With no signs of an AGN, the gas motion is most likely driven by the ongoing star formation. | With no signs of an AGN, the gas motion is most likely driven by the ongoing star formation. |
In both group members where we detect extraplanar ionized gas, the emission lines vary in terms of relative velocity and width from spaxel to spaxel indicating that there is no PSF broadening in these sources. | In both group members where we detect extraplanar ionized gas, the emission lines vary in terms of relative velocity and width from spaxel to spaxel indicating that there is no PSF broadening in these sources. |
As with $G1120-S1, we measure only motion along the line-of-sight while the gas is likely to be primarily moving perpendicular to the disk, i.e. the true gas velocities are likely to be higher. | As with SG1120-S1, we measure only motion along the line-of-sight while the gas is likely to be primarily moving perpendicular to the disk, i.e. the true gas velocities are likely to be higher. |
We cannot determine a net flow direction for the extraplanar gas because the errors on the systemic velocity (~100kms !) for these two galaxies are large compared to the velocity shifts (~10—65kms 1) in their extraplanar spaxels. | We cannot determine a net flow direction for the extraplanar gas because the errors on the systemic velocity $\sim100~\kms$ ) for these two galaxies are large compared to the velocity shifts $\sim 10-65 \kms$ ) in their extraplanar spaxels. |
However, we do confirm the existence of ionized gas at large scale heights above the disk of both members. | However, we do confirm the existence of ionized gas at large scale heights above the disk of both members. |
To determine what happens to the gas in these three members, we first estimate how much ionized gas is in the observed outflow. | To determine what happens to the gas in these three members, we first estimate how much ionized gas is in the observed outflow. |
For SG1120-S1, using the H5 lines from our single-slit data and the relation inoutflow.,, we assume case B recombination and an electron of 100cm? to estimate a total ionized gas mass of My;~10° in the two components of the HG line (Lag=2Mox1099erg s~!). | For SG1120-S1, using the $\beta$ lines from our single-slit data and the relation in, we assume case B recombination and an electron of $100 \cmc$ to estimate a total ionized gas mass of $M_{\mbox{\tiny
HII}}\sim 10^5\ \msol$ in the two components of the $\beta$ line $L_{H\beta}
= 2\times 10^{39}\ \ergsec$ ). |
Next we estimate an outflow rate (M) for the ionized gas by comparing the mass inferred from the Hj emission to a dynamical timescale. | Next we estimate an outflow rate $\dot{M}$ ) for the ionized gas by comparing the mass inferred from the $\beta$ emission to a dynamical timescale. |
Using the single- data we assume a radius of 1 kpc, consistent with the extent of the emission lines, and an outflow velocity of 900kms! from the most blueshifted component on the [OIIIJA5007 line, giving us tayn—R/V~10° yr. | Using the single-slit data we assume a radius of 1 kpc, consistent with the extent of the emission lines, and an outflow velocity of $900 \kms$ from the most blueshifted component on the $\lambda$ 5007 line, giving us $t_{dyn}=R/V\sim10^6$ yr. |
For | For |
fraction and angle. respectively. with We have applied this polarisation analysis to. our observations of the 620.701 GHz H»O 53».44, ortho-transition. | fraction and angle, respectively, with We have applied this polarisation analysis to our observations of the 620.701 GHz $\mathrm{H}_{2}\mathrm{O}$ $5_{32}-4_{41}$ ortho-transition. |
Our results shown in Figure | were obtained from the two aforementioned observations made at epochs of corresponding position angles of 261.27" and 277.46", | Our results shown in Figure \ref{fig:spectra} were obtained from the two aforementioned observations made at epochs of corresponding position angles of $261.27^{\circ}$ and $277.46^{\circ}$. |
Although there are no obvious strong polarisation signals from the maser emission peaks. we clearly detect polarisation levels ranging from pz1.5% to pz6% in regions of significant line intensity (e.. from approximately -5 to 45 km s7!). | Although there are no obvious strong polarisation signals from the maser emission peaks, we clearly detect polarisation levels ranging from $p\simeq1.5\%$ to $p\simeq 6\%$ in regions of significant line intensity (i.e., from approximately -5 to 45 km $^{-1}$ ). |
Furthermore. the observed anti-correlation of the polarisation fraction with the Stokes 7 intensity is similar to previous ground-based polarisation observations (Girartetal.. aimed at the detection of the Goldreich-Kylafis effect 1n. non-masing molecular lines (Goldreich&Kylafis.1981:Cortesetal.2005). which appears to have first been detected in evolved stars (Glennetal.. 1997). | Furthermore, the observed anti-correlation of the polarisation fraction with the Stokes $I$ intensity is similar to previous ground-based polarisation observations \citep{Girart2004,Hezareh2010} aimed at the detection of the Goldreich-Kylafis effect in non-masing molecular lines \citep{GK1981,Cortes2005}, which appears to have first been detected in evolved stars \citep{Glenn1997}. |
. We will discuss the relevance of the Goldreich-Kylatis effect for our observations in section 7 below. | We will discuss the relevance of the Goldreich-Kylafis effect for our observations in section 7 below. |
Several instrumental capabilities of HIFI and the observations obtained with them may be noted: (i) The two orthogonally polarised HIFI. receivers. are well matched and extremely stable. | Several instrumental capabilities of HIFI and the observations obtained with them may be noted: (i) The two orthogonally polarised HIFI receivers are well matched and extremely stable. |
But observations. of extended sources need to be conducted with caution. | But observations of extended sources need to be conducted with caution. |
A slight misalignment of H and V receivers can lead to a "false polarisation" that reverses polarity at half-year intervals (see Appendix Appendix A:)) ( | A slight misalignment of H and V receivers can lead to a “false polarisation" that reverses polarity at half-year intervals (see Appendix \ref{sec:1557}) ). ( |
11) The misalignment of the HIFI receivers does not appear to affect observations of unresolved sources. | ii) The misalignment of the HIFI receivers does not appear to affect observations of unresolved sources. |
Our observations realised at two observing epochs indicate that instrumental polarisation. which could be in part due to errors in the relative calibration between the two receiver chains of Band 1B. cannot exceed a measure of order |2% (see section 7 below). ( | Our observations realised at two observing epochs indicate that instrumental polarisation, which could be in part due to errors in the relative calibration between the two receiver chains of Band 1B, cannot exceed a measure of order $1-2\%$ (see section 7 below). ( |
it) The polarisation of VY CMais not significant near the peak of the 620.701 GHz H2O 53».44, line. but rises up to ~6% in the wings of the spectrum in a manner consistent with polarisation due to the Goldreich-Kylafis effect discussed in greater detail in section 7 below. ( | iii) The polarisation of VY CMa is not significant near the peak of the 620.701 GHz $\mathrm{H}_{2}\mathrm{O}$ $5_{32}-4_{41}$ line, but rises up to $\sim 6\%$ in the wings of the spectrum in a manner consistent with polarisation due to the Goldreich-Kylafis effect discussed in greater detail in section 7 below. ( |
1v) The stability of the 620.701 GHz masers 15 remarkable. | iv) The stability of the 620.701 GHz masers is remarkable. |
The variation over a three week periodis profile<1% (see Figure 2). ( | The variation over a three week period is $\lesssim1\%$ (see Figure 2). ( |
v) As Figure3. shows. the spectral of the 620.701 GHz and 22.235 GHz masers appears remarkably similar. | v) As Figure \ref{fig:22vs621} shows, the spectral profile of the 620.701 GHz and 22.235 GHz masers appears remarkably similar, |
The number of known astronomical sources of very-high-energy (VHE) y-rays grew ten-fold over the last five. | The number of known astronomical sources of very-high-energy (VHE) s grew ten-fold over the last five. |
. A large part of the newly discovered sources lie in the Galaxy and were revealed via a systematic scan of the inner Galactic Plane by the HESS telescope (Aharonianetal..2005a.2006). | A large part of the newly discovered sources lie in the Galaxy and were revealed via a systematic scan of the inner Galactic Plane by the HESS telescope \citep{HESS_survey_science,HESS_survey}. |
The HESS survey has covered an area 0.1 srin a strip ή<307. |b]<37. | The HESS survey has covered an area 0.1 sr in a strip $|l|<30^\circ$, $|b|<3^\circ$. |
This covers less than of the sky. | This covers less than of the sky. |
Surveys of larger regions on the VHE ssky with the existing ground based Cherenkov ttelescopes are difficult because the size of the field of view Is too narrow (5° for HESS. 3.5° VERITAS telescopes and 3° for MAGIC telescope). | Surveys of larger regions on the VHE sky with the existing ground based Cherenkov telescopes are difficult because the size of the field of view is too narrow $5^\circ$ for HESS, $3.5^\circ$ VERITAS telescopes and $3^\circ$ for MAGIC telescope). |
A previous survey of the northern hemisphere using the Cherenkov telescope Whipple has resulted only in derivation of upper limits on the flux of persistent VHE ssources (Weekesetal. 1979). | A previous survey of the northern hemisphere using the Cherenkov telescope Whipple has resulted only in derivation of upper limits on the flux of persistent VHE sources \citep{whipple_survey}. . |
The wide field of view MILAGRO (Atkinsetal..2004) and Tibet (Amenomonrtetal..2005) arrays have produced a systematic survey of the VHE ssky. | The wide field of view MILAGRO \citep{milagro} and Tibet \citep{tibet} arrays have produced a systematic survey of the VHE sky. |
However. the energy threshold of the air shower arrays like MILAGRO and TIBET is rather high (in the multi-TeV band) so that only sources with spectra extending well above | TeV could be detected. | However, the energy threshold of the air shower arrays like MILAGRO and TIBET is rather high (in the multi-TeV band) so that only sources with spectra extending well above 1 TeV could be detected. |
Contrary to the ground-based Cherenkov ttelescopes.Fermi has a wide field-of-view and continuously surveys the whole sky on a timescale of 3.2 hr. | Contrary to the ground-based Cherenkov telescopes, has a wide field-of-view and continuously surveys the whole sky on a timescale of $3.2$ hr. |
Over the first year of operation has detected some 1.5x10? Galactic and extragalactic sources of y-rays with energies above | GeV (Abdoetal..2009).. | Over the first year of operation has detected some $1.5\times 10^{3}$ Galactic and extragalactic sources of s with energies above 1 GeV \citep{fermi_catalog}. |
The smaller collection area of C Lim’. compared to ~10? n? for the ground-based ttelescopes) prevents an extension of the all-sky monitoring with to the VHE bband. | The smaller collection area of $\sim 1$ $^2$, compared to $\sim 10^5$ $^2$ for the ground-based telescopes) prevents an extension of the all-sky monitoring with to the VHE band. |
However. the collection area of is still sutficient for detecting the brightest ssources at the energies above 100 GeV. The power of the all- monitoring capabilities of at the highest energies was clearly demonstrated by discoveries of new VHE ssources motivated by detections of these sources above 10 GeV (Ong.2009.2010). | However, the collection area of is still sufficient for detecting the brightest sources at the energies above 100 GeV. The power of the all-sky monitoring capabilities of at the highest energies was clearly demonstrated by discoveries of new VHE sources motivated by detections of these sources above 10 GeV \citep{ATEL2260,ATEL2486}. |
. The all-sky survey capabilities of space-based ttelescope EGRET at the energies above 10 GeV were used for the search of new VHE bblazars by Dingus&Bertsch(2001):Gorbunovetal.(2005) via cross-correlation of arrival directions of highest energy EGRET photons with positions of known sources. | The all-sky survey capabilities of space-based telescope EGRET at the energies above 10 GeV were used for the search of new VHE blazars by \citet{dingus01,10GeV_EGRET} via cross-correlation of arrival directions of highest energy EGRET photons with positions of known sources. |
Below we use data to produce a survey of extragalactic sky at the energies above 100 GeV. re. in the energy range accessible for the ground-based ttelescopes. | Below we use data to produce a survey of extragalactic sky at the energies above 100 GeV, i.e. in the energy range accessible for the ground-based telescopes. |
We find that most of the sources visible with at the energies above 100 GeV are known TeV blazars. | We find that most of the sources visible with at the energies above 100 GeV are known TeV blazars. |
The only source which has not previously been reported as a VHE ssource turns out to be IC 310. which is a head-tail radio galaxy (Sibring&deBruyn.1998). with. possibly a BL Lae type nucleus (Rectoretal..1999). | The only source which has not previously been reported as a VHE source turns out to be IC 310, which is a head-tail radio galaxy \citep{sijbring98} with possibly a BL Lac type nucleus \citep{rector}. |
Two radio galaxies have previously been reported to be the sources of y-rays with energies above 100 GeV: M87 (Aharonianetal..2005b:Albert2008:Acciari2009) and Cen A (Aharonianetal..2009).. These two sources are the two closest Fanaroff-Riley type | (FR D) radio galaxies. | Two radio galaxies have previously been reported to be the sources of s with energies above 100 GeV: M87 \citep{m87,m87_magic,m87_veritas} and Cen A \citep{cena}.. These two sources are the two closest Fanaroff-Riley type I (FR I) radio galaxies. |
The FR I radio galaxies form the "parent population of BL Lac type blazars (Urry&Padovani.1995). | The FR I radio galaxies form the "parent" population of BL Lac type blazars \citep{urry95}. |
. They are expected to be weak VHE eemitters. because the ffülux from these sources is not boosted by the relativistic Doppler effect. | They are expected to be weak VHE emitters, because the flux from these sources is not boosted by the relativistic Doppler effect. |
Subsets and Splits
No saved queries yet
Save your SQL queries to embed, download, and access them later. Queries will appear here once saved.